diff --git "a/batch_s000026.csv" "b/batch_s000026.csv" new file mode 100644--- /dev/null +++ "b/batch_s000026.csv" @@ -0,0 +1,10368 @@ +source,target + Despite much effort. little is presently known about the nature of these sources. particularly since the counterparts at. other wavelengths are difficult to find.," Despite much effort, little is presently known about the nature of these sources, particularly since the counterparts at other wavelengths are difficult to find." + Indeed. these sources are often superposed against regions of high surface brightness in their host galaxy.," Indeed, these sources are often superposed against regions of high surface brightness in their host galaxy." + Thus. most of the optical identifications are obtained by spatial coincidence of the counterparts and comparison with broadband spectral characteristics of known stars in the host (Liu et al.," Thus, most of the optical identifications are obtained by spatial coincidence of the counterparts and comparison with broadband spectral characteristics of known stars in the host (Liu et al." + 2002; Wu et al., 2002; Wu et al. + 2002: Zezas et al., 2002; Zezas et al. + 2002: Zampieri et al., 2002; Zampieri et al. + 2003)., 2003). + In other cases. it has been possible to study only the nearby environment in the ULX host (Pakull Mirtont 2002: Wang 2002: Roberts et al.," In other cases, it has been possible to study only the nearby environment in the ULX host (Pakull Mirioni 2002; Wang 2002; Roberts et al." + 2003)., 2003). + In one case (NGC4698-ULX1) the optical spectral features allowed a clear dentification as a background BL Lae object (Foschini et al., In one case (NGC4698-ULX1) the optical spectral features allowed a clear identification as a background BL Lac object (Foschini et al. + 20022)., 2002a). + Here we report the identification of the nature of another ULX., Here we report the identification of the nature of another ULX. + The counterparts in the infrared and optical bands were found. and spectroscopy revealed clear Balmer and forbidden transition emission lines.," The counterparts in the infrared and optical bands were found, and spectroscopy revealed clear Balmer and forbidden transition emission lines." + The derived redshift of =0.217 indicates a background galaxy in this case also., The derived redshift of $z=0.217$ indicates a background galaxy in this case also. + GC 4168 ts an E2 elliptical galaxy located in the Virgo cluster (d=16.8 Mpc)., NGC 4168 is an E2 elliptical galaxy located in the Virgo cluster $d=16.8$ Mpc). + It hosts an active galactic nucleus (AGN). classified as a Seyfert 1.9 by Ho et al. (," It hosts an active galactic nucleus (AGN), classified as a Seyfert 1.9 by Ho et al. (" +1997).,1997). + The galaxy was observed on 4 December 2001 using the European Photon Imaging Camera (EPIC) on board the satellite (see Foschini et al., The galaxy was observed on 4 December 2001 using the European Photon Imaging Camera (EPIC) on board the satellite (see Foschini et al. + 20020: for the X-ray part. we refer in the following to the results obtained in this paper. unless explicitly stated).," 2002b; for the X-ray part, we refer in the following to the results obtained in this paper, unless explicitly stated)." + EPIC is composed of two instruments: the PN-CCD camera (Strüdder et al., EPIC is composed of two instruments: the PN-CCD camera (Strüdder et al. + 2001) and two MOS-CCD detectors (Turner et al., 2001) and two MOS-CCD detectors (Turner et al. + 2001)., 2001). + The effective exposure time was 17.4 ks., The effective exposure time was $17.4$ ks. + One ULX was found apparently associated with NGC 4168. being inside the galaxy’s Dos ellipse. at 45” from the optical centre of the galaxy.," One ULX was found apparently associated with NGC 4168, being inside the galaxy's $D_{25}$ ellipse, at $45''$ from the optical centre of the galaxy." +" The ULX has coordinates (2000) à=12h|]2""[5 and 6=+13°12/48”. with an uncertainty radius of 4."," The ULX has coordinates (J2000) $\alpha = 12^{\rm h} 12^{\rm m} +14\fs5$ and $\delta = +13^{\circ} 12\arcmin 48\arcsec$, with an uncertainty radius of $4''$." + The counts were not sufficient to extract à spectrum. so we could only convert to physical units using the count rates derived from the task of the XMM-SAS software (v. 5.2).," The counts were not sufficient to extract a spectrum, so we could only convert to physical units using the count rates derived from the task of the XMM-SAS software (v. 5.2)." +" We found a count rate of 4.0+0.7 counts s! in the 0.5-10 keV band. which corresponds to a flux of (1.8+0.3)x107 ere em7 s! adopting a conversion factor of 3x10"" counts env erg!. which in tum was derived by using a power-law model with [=2.0 and an average Galactic column density of Ny=3x10-° em7?."," We found a count rate of $4.0 \pm 0.7$ counts $^{-1}$ in the 0.5–10 keV band, which corresponds to a flux of $(1.8\pm 0.3)\times +10^{-14}$ erg $^{-2}$ $^{-1}$ adopting a conversion factor of $3\times 10^{11}$ counts $^2$ $^{-1}$, which in turn was derived by using a power-law model with $\Gamma=2.0$ and an average Galactic column density of $N_{\rm H} = 3\times 10^{20}$ $^{-2}$." + This value for Tis common among the ULXs found withXMM-Newton., This value for $\Gamma$ is common among the ULXs found with. +. This flux value was corrected according to the energy encircled fraction (Ghizzardi 2001)., This flux value was corrected according to the energy encircled fraction (Ghizzardi 2001). + Correction for vignetting (Lumb 2002) was not applied becausethe source is close to the center of the field of view (<2’)., Correction for vignetting (Lumb 2002) was not applied becausethe source is close to the center of the field of view $<$$2\arcmin$ ). + At the distance of NGC 4168 the resulting 0.5—10 keV luminosity is 6»1075 erg s! (for a discussion on the luminosity threshold of ULXs. see Foschini et al.," At the distance of NGC 4168 the resulting $0.5-10$ keV luminosity is $6\times 10^{38}$ erg $^{-1}$ (for a discussion on the luminosity threshold of ULXs, see Foschini et al." + 2002c)., 2002c). + A single optical object with coordinates. (J2000) 12{2111460 and 6=137 1247/8 (i.e.. [4 from the position. thus well within the X-ray error box) is clearly visible on the Digitized Sky Survey observation of NGC 4168. originally made with the 48-inch Schmidt telescope at Palomar Observatory on 14 April 1955.," A single optical object with coordinates (J2000) $\alpha = 12^{\rm h} +12^{\rm m} 14\fs60$ and $\delta = +13^{\circ} 12\arcmin 47\farcs$ 8 (i.e., $\farcs$ 4 from the position, thus well within the X-ray error box) is clearly visible on the Digitized Sky Survey observation of NGC 4168, originally made with the 48-inch Schmidt telescope at Palomar Observatory on 14 April 1955." + We thus consider it as the optical counterpart of NGC4168-ULX1., We thus consider it as the optical counterpart of NGC4168-ULX1. + The source is also present in the US Naval Observatory (USNO) BI.O Catalog (Monet et al., The source is also present in the US Naval Observatory (USNO) B1.0 Catalog (Monet et al. + 2003) with the identification number 1032—0222128., 2003) with the identification number $1032-0222128$. + The magnitudes in the different bands are BI=19.1. B2=17.8. RI=18.1. R2=18.2. and J=17.8.," The magnitudes in the different bands are $B1=19.1$, $B2=17.8$, $R1=18.1$, $R2=18.2$, and $I=17.8$." + The BI and RI magnitudes refer to the Palomar Observatory Sky Survey I (POSSI). performed between 1949 and 1965.," The $B1$ and $R1$ magnitudes refer to the Palomar Observatory Sky Survey I (POSSI), performed between 1949 and 1965." + The B2. R2 and / magnitudes are measured from the Palomar Observatory Sky Survey II (POSSID. performed from 1985 to 2000.," The $B2$ , $R2$ and $I$ magnitudes are measured from the Palomar Observatory Sky Survey II (POSSII), performed from 1985 to 2000." +"timescale near the WD photosphlere. aud the Iuuinosity ofthe svstei is dominateds by the WD's quiesceut surface Iuniuositv. L,. which originates froii deep boucath the photosphere (Piroetal.2005).","timescale near the WD's photosphere, and the luminosity of the system is dominated by the WD's quiescent surface luminosity, $L_q$, which originates from deep beneath the photosphere \citep{piro05}." +. Ever since Sion(1990) diseussed the value of Tig measurements of dwarf novae iu quiescence. observers and theorists have vigorously pursued this important diagnostic.," Ever since \cite{sion99} discussed the value of $T_{\rm +eff}$ measurements of dwarf novae in quiescence, observers and theorists have vigorously pursued this important diagnostic." + Most receutlv. Townsley&Gausicke(2009) sunuuarized the observations aud improved the earlier theoretical work (Townsley&Bildsten2002: 2005)..," Most recently, \cite{tg09} summarized the observations and improved the earlier theoretical work \citep{tb02,godon02,tb03,tb05,piro05}. ." + In Figure 8.. we have added the Tuy evolutions for our two scenarios to the original figure from Townsley&Cansicke (2009).," In Figure \ref{fig:teffevol}, we have added the $T_{\rm eff}$ evolutions for our two scenarios to the original figure from \cite{tg09}." +. Some data are marginally cousisteut with the solid line (scenario A). whereas no observed WDs are as cold as predicted from our scenario D: however. it is unclear if the abseuce of these cold systems is physical or due to selection effects;," Some data are marginally consistent with the solid line (scenario A), whereas no observed WDs are as cold as predicted from our scenario B; however, it is unclear if the absence of these cold systems is physical or due to selection effects." + It would be of iuterest to learn whether those cold svstenis near the solid line have anv other evidence for having low-1ass accretors., It would be of interest to learn whether those cold systems near the solid line have any other evidence for having low-mass accretors. + The coutinnal discoveries of CVs in quiescence by the Sloan Digital Sky Survev (e.g. Szkodvetal.20001: Gansickeetal. 20093) and. im the near future. bv SkvMapper (Murphyctal.2008) will certainly provide new opportunitics to reveal acercting Πο WDs.," The continual discoveries of CVs in quiescence by the Sloan Digital Sky Survey (e.g., \citealt{szkody09}; \citealt{gaen09}) ) and, in the near future, by SkyMapper \citep{skymapper} will certainly provide new opportunities to reveal accreting He WDs." + Iu particular. Causickeetal. (200937s recent discovery of the expected “pile-up” of CVs at the SOδ min orbital period miumnmn has alleviated long-standing— coucerus reearding binary evolution.," In particular, \cite{gaen09}' 's recent discovery of the expected “pile-up” of CVs at the $80-86$ min orbital period minimum has alleviated long-standing concerns regarding binary evolution." + Heuce. as we noted iu the introduction. z20% of the CVs in Cansickeetal. (2009) compilation should harbor a Πο WD.," Hence, as we noted in the introduction, $\approx 20\%$ of the CVs in \cite{gaen09}' 's compilation should harbor a He WD." + The discoverys and study of accreting WD. pulsators (see Mikadaimetal.2007 for an updated list) may well be our best hope. as Arrasetal.(2006) noted that low-1ass WDs have colder effective teirperatures for pulsation than higher mass WDs. aud our Z;4 calculations shown in Fieure δ are sliehtlv colder than the blue edge calculated by Arrasetal.(2006) for low eravity (6.9. low-uass WDz).," The discovery and study of accreting WD pulsators (see \citealt{muka07} for an updated list) may well be our best hope, as \cite{arras06} noted that low-mass WDs have colder effective temperatures for pulsation than higher mass WDs, and our $T_{\rm eff}$ calculations shown in Figure \ref{fig:teffevol} are slightly colder than the blue edge calculated by \cite{arras06} for low gravity (e.g., low-mass WDs)." + Thus. a prevalence of pulsators at low Tig may be an iudicator of Te WDs.," Thus, a prevalence of pulsators at low $T_{\rm eff}$ may be an indicator of He WDs." + Certainly more work. is needed to make this connection clear. but the rapid increase in the discovery of such pulsators is bound to reveal a few new systems wortliv of iuteusive study.," Certainly more work is needed to make this connection clear, but the rapid increase in the discovery of such pulsators is bound to reveal a few new systems worthy of intensive study." + Iu the previous sections. we studied the secular evolution of the WD.," In the previous sections, we studied the secular evolution of the WD." + We now turn our attention to details of the individual nova outbursts. focusing onu the evolution of the convective burning phase. the composition of the nova ejecta. aud the WD's post-nova appearance as a supersoft X-ray source.," We now turn our attention to details of the individual nova outbursts, focusing on the evolution of the convective burning phase, the composition of the nova ejecta, and the WD's post-nova appearance as a supersoft x-ray source." + For most of the accretion phase of the nova cycle. he thermal profile is set by “compressional” heating (seo §3.2)).," For most of the accretion phase of the nova cycle, the thermal profile is set by “compressional” heating (see \ref{sec:constMdot}) )." + Ποπονα. ouce the base of the envelope vecolmcs dense aud hot enoush. the energy generation rate from nuclear burning becomes large enough that radiative diffusion cau uo longer effectively transport the uumositv.," However, once the base of the envelope becomes dense and hot enough, the energy generation rate from nuclear burning becomes large enough that radiative diffusion can no longer effectively transport the luminosity." + The radiative cuvelope is then transformed into a convective zone whose cutropy is increased * continued nuclear burning until mass is lost via a radiativelv-diiveu wind. aud the nova outburst is observed.," The radiative envelope is then transformed into a convective zone whose entropy is increased by continued nuclear burning until mass is lost via a radiatively-driven wind, and the nova outburst is observed." + We now give a brief overview of the rolevaut physics cing this stage of the CN evolution: for a more coniplete explanation of simular physics iu IHe-burniug convective envelopes. see 83 of Shen&Bildsten(200900).," We now give a brief overview of the relevant physics during this stage of the CN evolution; for a more complete explanation of similar physics in He-burning convective envelopes, see 3 of \cite{sb09b}." + Figure 9 shows au example of the evolution of a convective euvelope as it is heatedby nuclear burning. nunmercallv calculated for à 0.4AL. WD that accreted mass at a rate of 10t+AL.wv+ after it has reached its equilibrium core temperature of 5.5.10° K. Dotted," Figure \ref{fig:rhotzone} shows an example of the evolution of a convective envelope as it is heatedby nuclear burning, numerically calculated for a $0.4 \msol$ WD that accreted mass at a rate of $10^{-11} \smpy$ after it has reached its equilibrium core temperature of $5.5\E{6}$ K. Dotted" +The infrared radiation fro1l ΠΙΟΣ star-forming galaxies is dominated by eiission YOM ¢ust erains heated by absorbed stellay enerev.,The infrared radiation from most star-forming galaxies is dominated by emission from dust grains heated by absorbed stellar energy. + Dust enüssion is powered by absorption of radiation frou loniziug and non-ioniziug stars., Dust emission is powered by absorption of radiation from ionizing and non-ionizing stars. + Dust is nost efficient at absorbing photons in the ultraviolet. (UV) as the reative optical depth of dust is the highest in the UV (Cordonetal.2003)., Dust is most efficient at absorbing photons in the ultraviolet (UV) as the relative optical depth of dust is the highest in the UV \citep{Gordon2003}. +.. Oulv carly type CO and D) stars produce sienificaut amounts of UV photons: however. frese hot massive stars have short lifetimes (less than 100 nüllion vears) ancl are formed in relatively small uunbers compared to less massive. less luminous. auxL cooler stars that produce very few UV photons.," Only early type (O and B) stars produce significant amounts of UV photons; however, these hot massive stars have short lifetimes (less than 100 million years) and are formed in relatively small numbers compared to less massive, less luminous, and cooler stars that produce very few UV photons." + Caven the initial mass ftnction and evolutionary history. this implies that star-forming ealaxies have a sinal lass fraction of UW bright. vouug stars as compared o UV faint. old stars.," Given the initial mass function and evolutionary history, this implies that star-forming galaxies have a small mass fraction of UV bright, young stars as compared to UV faint, old stars." + Thus. the question arises: whic1 population of stars coluinates the dust heating in star-forming clouds?," Thus, the question arises: which population of stars dominates the dust heating in star-forming clouds?" + The less numerous but much brighter iu the UN “young stars («100 Myr) or the unuuerous but uich faiuter iu the UV od stars?," The less numerous but much brighter in the UV young stars $< +100$ Myr) or the numerous but much fainter in the UV old stars?" + Tow does this answer change when we consider the ciission at specific IR wavelengths?, How does this answer change when we consider the emission at specific IR wavelengths? + The majority of the IR energv from star-forming ealaxies is endtted at £u-IR 100 μπι) wavelengths., The majority of the IR energy from star-forming galaxies is emitted at far-IR 100 ) wavelengths. + Tistorically. this fav-IR cuuission has been identified as intrared crus emission from dust heated by nou-donizig populations (IIelou199L).. which ave older than 10 My.," Historically, this far-IR emission has been identified as infrared cirrus emission from dust heated by non-ionizing populations \citep{Helou1994}, which are older than 10 Myr." + LousdalePersson&IIlou(1987) interpreted the far-IR cussion from spiral disks iu terms of two thermal COMPOleits with different temperatures and found tha the cirus componcut contributes amore than half of the total far-IR fiux., \citet{Cirrus1987} interpreted the far-IR emission from spiral disks in terms of two thermal components with different temperatures and found that the cirrus component contributes more than half of the total far-IR flux. + However. as old stars cuit very few of the UV. photons hat power dust enüssion. 1 is possibe that UV-brigh voung stars could dominate cirus eniüsson.," However, as old stars emit very few of the UV photons that power dust emission, it is possible that UV-bright young stars could dominate cirrus emission." + For exame. a stnall uuuber of vouug stars embedded ina larexορ optically thin cloud cau resul in a dihte radiation field. cold dust temperature and therefore cold cirrus emission.," For example, a small number of young stars embedded in a large optically thin cloud can result in a dilute radiation field, cold dust temperature and therefore cold cirrus emission." + Star formation rate (SER) indicators are im)ortiui observational probes of the star formation histories of ealaxies., Star formation rate (SFR) indicators are important observational probes of the star formation histories of galaxies. + They are ustally sinele-band or waveleugsth-integrated. quantities hat are presumed to trace a specific reguue of recaut star formation du a region or galaxv (Ixeunicutt1998)., They are usually single-band or wavelength-integrated quantities that are presumed to trace a specific regime of recent star formation in a region or galaxy \citep{Kennicutt1998}. + The most conunon SER indicators include the Πα fiux (tracing unocured loniziug stars. $ 3.9 ), $>$ 3.8 ), and $>$ 3.9 ) milli-mags over 1 hour time-scales" +(centred at 23.75 Watt/I1z) containing only a single source (BOS830|5813) is omitted clue to its large uncertainty.,(centred at 23.75 Watt/Hz) containing only a single source (B0830+5813) is omitted due to its large uncertainty. +" We compared the resulting LLE with an LLP of a simulated: radio source population of 10"" objects. with random ages. and a jet-power distribution defined as in equation LO."," We compared the resulting LLF with an LLF of a simulated radio source population of $10^6$ objects, with random ages, and a jet-power distribution defined as in equation 10." +" The ""observed? Luminosity of à source was calculated assuming that it hack evolved. over its lifetime according to the luminosity evolution derived in section 4.1. out to a maximum size. r ."," The `observed' luminosity of a source was calculated assuming that it had evolved over its lifetime according to the luminosity evolution derived in section 4.1, out to a maximum size, $r_+$ ." +" Mr«ors. the size of the source evolves as r=(£2PET,"," At $rr_*$." + This results in a similar evolution for large size radio sources as derived in section 4.1. with the luminosity aber=gu only dependent on £27.," This results in a similar evolution for large size radio sources as derived in section 4.1, with the luminosity at $r=r_*$ only dependent on $P_J$." + Lo was not our aim to determine absolute values for number densities ancl racio powers with these simulations., It was not our aim to determine absolute values for number densities and radio powers with these simulations. + The results of the simulation were scaled in such way. that the LLE obtained for large size radio sources. matched the LLP of steep spectrum. sources as derived by Dunlop peacock (1990).," The results of the simulation were scaled in such way, that the LLF obtained for large size radio sources, matched the LLF of steep spectrum sources as derived by Dunlop peacock (1990)." + ‘Table 4 lists the important characteristics of the simulated LLE of large size and GPS sources. and their dependence on the free model parameters.," Table \ref{param} lists the important characteristics of the simulated LLF of large size and GPS sources, and their dependence on the free model parameters." + The parameters ó and 2. as defined in equations LO and 3. determine the slope of the low ancl high luminosity part of the LLP of aree size radio sources.," The parameters $\delta$ and $\beta$, as defined in equations 10 and 3, determine the slope of the low and high luminosity part of the LLF of large size radio sources." + These were chosen to be similar to he parameters @1l and ὁ1 as derived by Dunlop Peacock (1990). with &=1.10 and 3=1.16.," These were chosen to be similar to the parameters $a-1$ and $b-1$ as derived by Dunlop Peacock (1990), with $\delta=-1.10$ and $\beta=1.16$." + This value of ds slightly lower than derived (rom X-ray observations of nearby ellipticals (3=1.52. Trinchieri et al.," This value of $\beta$ is slightly lower than derived from X-ray observations of nearby ellipticals $\beta=1.5-2$, Trinchieri et al." + 1986)., 1986). + ote however. that the radio source population is dominated » objects with size z20 kpe. for which the surrounding medium is dominated by intra-cluster gas. which is expected o have a [Blatter density. gradient.," Note however, that the radio source population is dominated by objects with size $>20$ kpc, for which the surrounding medium is dominated by intra-cluster gas, which is expected to have a flatter density gradient." + With the parameters ó and 3 and the slopes of the LLP of large size. radio sources fixed. the relative positions of the break luminosities could. be determined.," With the parameters $\delta$ and $\beta$ and the slopes of the LLF of large size radio sources fixed, the relative positions of the break luminosities could be determined." +" A sharp cut-olf will occur near the highest luminosity. L,,,,."," A sharp cut-off will occur near the highest luminosity, $L_{max}$." + Phe number of GPS galaxies in the highest luminosity bin. as shown in figure SN.. is lower han expected from the extrapolation of the LLE at lower uminosities.," The number of GPS galaxies in the highest luminosity bin, as shown in figure \ref{llf}, , is lower than expected from the extrapolation of the LLF at lower luminosities." + Εις can be explained if this luminosity bin is near the cut-oll. luminosity Lue., This can be explained if this luminosity bin is near the cut-off luminosity $L_{max}$ . + We therefore. chose og Linas to be 27.1 (W 1)., We therefore chose log $L_{max}$ to be 27.1 (W $^-1$ ). + Phe break luminosity of arge size radio sources. is also determined by Dunlop Peacock (1990) to be log Lys. = 25.79 (corrected. to 5 (01414).," The break luminosity of large size radio sources, is also determined by Dunlop Peacock (1990) to be log $L_{LS*}$ = 25.79 (corrected to 5 GHz)." + As can be seen from table 4.. the luminosity ratio LancefLis determines the value of &fry.," As can be seen from table \ref{param}, the luminosity ratio $L_{max}/L_{ls*}$ determines the value of $r_{+}/r_{*}$." + Fhis corresponds o à maximum size for a radio source of LOO kpe. assuming ry=| kpc.," This corresponds to a maximum size for a radio source of 100 kpc, assuming $r_*=1$ kpc." + This value is quite near the turnover seen in the linear size distribution of 3€1t. galaxies. as shown w O'Dea Baum (1997).," This value is quite near the turnover seen in the linear size distribution of 3CR galaxies, as shown by O'Dea Baum (1997)." + The break luminosity of GPS sources. Lypss. relative to Line. is dependent on the range of jet-powers CP/P.).," The break luminosity of GPS sources, $L_{gps*}$, relative to $L_{max}$, is dependent on the range of jet-powers $(P_+/P_-)$." + To let Lypss coincide with the peak in the observed. GPS LLE. a value of (2οJ=200 was used.," To let $L_{gps*}$ coincide with the peak in the observed GPS LLF, a value of $(P_+/P_-)$ =200 was used." + Although the uncertainties on the datapoints are large ancl several free. parameters enter the simulation. figure shows that the shape of the LLE of GPS sources is as expected.," Although the uncertainties on the datapoints are large and several free parameters enter the simulation, figure \ref{llf} shows that the shape of the LLF of GPS sources is as expected." + Note that most free parameters are determined by fitting the LLE of large size radio sources to that of Dunlop Peacock (1990). except rand (P./P.2.," Note that most free parameters are determined by fitting the LLF of large size radio sources to that of Dunlop Peacock (1990), except $r_{+}$ and $(P_+/P_-)$." + This analysis should. be regarded as an example of how future. large ancl homogeneously defined: samples of GPS sources can constrain the luminosity evolution. of extragalactic racio SOULCCS., This analysis should be regarded as an example of how future large and homogeneously defined samples of GPS sources can constrain the luminosity evolution of extragalactic radio sources. + The proposed. increase in luminosity for voung racio sources seems to be in contradiction to the high number counts ofGPS sources with respect to large size racio sources suggesting that they should. decrease in radio Luminosity. o» a factor 10 during their. lifetime (Fanti ct al., The proposed increase in luminosity for young radio sources seems to be in contradiction to the high number counts of GPS sources with respect to large size radio sources suggesting that they should decrease in radio luminosity by a factor $~10$ during their lifetime (Fanti et al. + 1995. veadhbead ct al.," 1995, Readhead et al." + 1996. οDea Daum 1997).," 1996, O'Dea Baum 1997)." + Lowever. this is not the case.," However, this is not the case." + Flux density. limited samples. as used. for hese analyses. only. probe the most luminous objects at any redshift.," Flux density limited samples, as used for these analyses, only probe the most luminous objects at any redshift." + As can be seen from figure S.. at high luminosities. he two luminosity function approach each other. due to the latter slope of the Iuminosity function of GPS sources.," As can be seen from figure \ref{llf}, at high luminosities, the two luminosity function approach each other, due to the flatter slope of the luminosity function of GPS sources." + This results in a relatively high number density of GPS sources in flux density limited samples., This results in a relatively high number density of GPS sources in flux density limited samples. + In this paper we show that in addition to the well known correlation between spectral peak frequeney. and angular size (eg., In this paper we show that in addition to the well known correlation between spectral peak frequency and angular size (eg. + Fanti ct al., Fanti et al. + 1990). a correlation exists between the peak flux density and angular size of GPS CSS sources.," 1990), a correlation exists between the peak flux density and angular size of GPS CSS sources." + The strength. anc sign of these correlations are exactly as expected from SSA theory. assuming equipartition. and are therefore a strong indication that SSA is indeed the cause of the spectral turnovers in these objects.," The strength and sign of these correlations are exactly as expected from SSA theory, assuming equipartition, and are therefore a strong indication that SSA is indeed the cause of the spectral turnovers in these objects." + Furthermore. these correlations are consistent with GPSCSS sources evolvingin a self-similar way.," Furthermore, these correlations are consistent with GPSCSS sources evolvingin a self-similar way." + Interestingly. the scll-similar evolution scenario is better fitted by assuming an equipartition than a constant magneticfield.," Interestingly, the self-similar evolution scenario is better fitted by assuming an equipartition than a constant magneticfield." + In Duxdensity limited samples. GPS galaxies are found at higher redshifts than large size radio sources.," In fluxdensity limited samples, GPS galaxies are found at higher redshifts than large size radio sources." + Since the, Since the +GRBO70616 is intriguing in that the emission rises relatively slowly over 100 seconds to a peak. then persists at a fairly constant level before showing a very rapid decline.,"GRB070616 is intriguing in that the emission rises relatively slowly over 100 seconds to a peak, then persists at a fairly constant level before showing a very rapid decline." + GRB 070129 is similar to GRB 070721B in that it has a possible internal plateau that is interrupted by a flare followed by a steep decline., GRB 070129 is similar to GRB 070721B in that it has a possible internal plateau that is interrupted by a flare followed by a steep decline. + GRB 070110 displays a canonical early light curve with an initial steep decline. but then exhibits a period of relatively constant emission.," GRB 070110 displays a canonical early light curve with an initial steep decline, but then exhibits a period of relatively constant emission." + Following this plateau the decay is surprisingly steep a~ 7) decay (Troja 22007)., Following this plateau the decay is surprisingly steep $\alpha \sim 7$ ) decay (Troja 2007). + Thus in this case the proto-magnetar survived for much longer than in most of the other GRBs., Thus in this case the proto-magnetar survived for much longer than in most of the other GRBs. + GRB 060607A appears to follow the canonical lighteurve with a “normal” X-ray plateau with multiple flares preventing a good fit with the two component model.," GRB 060607A appears to follow the canonical lightcurve with a ""normal"" X-ray plateau with multiple flares preventing a good fit with the two component model." + However at late times the decay following the plateau is too steep for an afterglow and is consistient with a ~4.," However at late times the decay following the plateau is too steep for an afterglow and is consistient with $\alpha$ $\sim +4$." + This is unlikely to be explained by anvthing other than central engine activity and thus has been included in the internal plateau sample., This is unlikely to be explained by anything other than central engine activity and thus has been included in the internal plateau sample. + As in GRB 070110. the internal plateau seen in GRB 060607A dominates the burst emission unusually late starting at about 900 seconds when (from Table 1) most of the other internal plateaus have ended.," As in GRB 070110, the internal plateau seen in GRB 060607A dominates the burst emission unusually late starting at about 900 seconds when (from Table 1) most of the other internal plateaus have ended." + GRB 060510B (also shown in Fig. Lis , GRB 060510B (also shown in Fig. \ref{fitexamples}) ) +very similar to GRB 070616., is very similar to GRB 070616. + In both cases the proposed internal plateau dominates the emission from the burst very early on., In both cases the proposed internal plateau dominates the emission from the burst very early on. + GRB 060202 displays unusual emission attributed to an internal plateau between 325 and 766 seconds., GRB 060202 displays unusual emission attributed to an internal plateau between 325 and 766 seconds. + The fluctuations during this plateau are unusually regular., The fluctuations during this plateau are unusually regular. + GRB 050904 has multiple flares at early and late times. but at about 230 seconds there is a period where the emission appears relatively constant followed by a steep decay. leading it to be included in the sample as a possible internal plateau.," GRB 050904 has multiple flares at early and late times, but at about 230 seconds there is a period where the emission appears relatively constant followed by a steep decay, leading it to be included in the sample as a possible internal plateau." + To further investigate the nature of the internal plateau we compared the X-ray data to optical/UV data from the UVOT., To further investigate the nature of the internal plateau we compared the X-ray data to optical/UV data from the UVOT. + The GRBs within the sample with near-simultaneous optical/UV and X-ray light curves are shown in Fig. 3.., The GRBs within the sample with near-simultaneous optical/UV and X-ray light curves are shown in Fig. \ref{optical2}. + While an early rise in the optical can be seen in some cases. the optical emission does not show the same behaviour as the X-ray.," While an early rise in the optical can be seen in some cases, the optical emission does not show the same behaviour as the X-ray." + The internal plateau and following steep decay are significantly more prominent in X-rays., The internal plateau and following steep decay are significantly more prominent in X-rays. + For example in GRB 070616. the optical is constant from before the plateau in the X-ray and until after the steep decline.," For example in GRB 070616, the optical is constant from before the plateau in the X-ray and until after the steep decline." + In Fig., In Fig. + 2 if the plateau seen in each of the X-ray lightcurves is of an external origin. then the X-ray and optical lighteurve should be related to each other in a manner consistent with the external shock model. i.e. the breaks should be achromatic.," \ref{sample} if the plateau seen in each of the X-ray lightcurves is of an external origin, then the X-ray and optical lightcurve should be related to each other in a manner consistent with the external shock model, i.e. the breaks should be achromatic." + However. if the X-ray and optical emission components are not related to each other. e.g. a sharp decay in X-ray but no break in optical. this strongly suggests that the X-ray emission is not external or a jet-break but rather is of internal origin.," However, if the X-ray and optical emission components are not related to each other, e.g. a sharp decay in X-ray but no break in optical, this strongly suggests that the X-ray emission is not external or a jet-break but rather is of internal origin." + In Troja (2007) for GRB 0701[0 four spectral energy distributions (SEDs) were examined during the initial decay. the beginning and end of the plateau and during the shallow decay after the steep decline.," In Troja (2007) for GRB 070110 four spectral energy distributions (SEDs) were examined during the initial decay, the beginning and end of the plateau and during the shallow decay after the steep decline." + These SEDs were constructed by extrapolating the X-ray spectrum to the lower energies., These SEDs were constructed by extrapolating the X-ray spectrum to the lower energies. + During the initial decay the optical data are not consistent with the extrapolation of the X-ray spectrum to low energies., During the initial decay the optical data are not consistent with the extrapolation of the X-ray spectrum to low energies. + During the internal plateau. the optical and X-ray spectral distributions are also completely inconsistent with one another. implying different origins for the optical and X-ray photons.," During the internal plateau, the optical and X-ray spectral distributions are also completely inconsistent with one another, implying different origins for the optical and X-ray photons." + For GRB 080310 and GRB 070616 the extrapolation of the X-ray spectrum is also inconsistient with the optical during the internal plateau (Beardmore iin preparation. Starling 22008).," For GRB 080310 and GRB 070616 the extrapolation of the X-ray spectrum is also inconsistient with the optical during the internal plateau (Beardmore in preparation, Starling 2008)." + Likewise. for GRBO60607A extrapolating the X-ray spectum to the optical in a similar way to Troja," Likewise, for GRB060607A extrapolating the X-ray spectum to the optical in a similar way to Troja" +The same-polarity encountering results in a merging of those elements to the flux related to sunspols.,The same-polarity encountering results in a merging of those elements to the flux related to sunspots. +" Whereas. (he opposite polarity encountering causes [Iux cancelations with the net results of lost smaller elements anc a diffusion of sunspot Πας,"," Whereas, the opposite polarity encountering causes flux cancelations with the net results of lost smaller elements and a diffusion of sunspot flux." + What accompanied the sunspol flux diffusion is the reduced smaller elements with the turbulent origin., What accompanied the sunspot flux diffusion is the reduced smaller elements with the turbulent origin. + This accounts [or the anti-correlated magnetic component possibly., This accounts for the anti-correlated magnetic component possibly. + By this kind of interaction magnetic flux from turbulent dvnamo actively takes part in the operation of the solar cvele. helping with more ellicient magnetic diffusion.," By this kind of interaction magnetic flux from turbulent dynamo actively takes part in the operation of the solar cycle, helping with more efficient magnetic diffusion." + To quantify Chis mechanism. studies of dynamic interaction between small-scale magnetic elements ancl active regions fields are crucially required.," To quantify this mechanism, studies of dynamic interaction between small-scale magnetic elements and active regions fields are crucially required." + secondly. itis also possible that at the solar maximum. the stronger magnetic field from sunspots tends to suppress the Sun'ss global convection in some measure.," Secondly, it is also possible that at the solar maximum, the stronger magnetic field from sunspots tends to suppress the s global convection in some measure." + As a result. the local dvnamo has been abated somehow. and the network elements created by turbulence are reduced in number and total Πας.," As a result, the local dynamo has been abated somehow, and the network elements created by turbulence are reduced in number and total flux." + This seems to suggest that the (turbulent dynamo is. in fact. global but not local.," This seems to suggest that the turbulent dynamo is, in fact, global but not local." + Unfortunately. so far there have been no definite observations about the changes in the global solar convection during the suuspot cycle.," Unfortunately, so far there have been no definite observations about the changes in the global solar convection during the sunspot cycle." + Another possibility is that the anti-correlated component represents the recveling of parts of the previously diffused or submerged magnetic flux from the mean-field dvnamo (Parker 1937)., Another possibility is that the anti-correlated component represents the recycling of parts of the previously diffused or submerged magnetic flux from the mean-field dynamo (Parker 1987). + The diffusion of magnetic flux [rom sunspots to the deep convection zone requires 5-7 vears (Jiang οἱ al., The diffusion of magnetic flux from sunspots to the deep convection zone requires 5-7 years (Jiang et al. + 2007)., 2007). + Parts of the cdiffised or submerged flux serves as the seed field [ον the elobally turbulent. dvuamo., Parts of the diffused or submerged flux serves as the seed field for the globally turbulent dynamo. + Its production is naturally out of phase with sunspols in the solar cvcle. and brings up the magnetic elements (hat. anti-phased. with sunspots.," Its production is naturally out of phase with sunspots in the solar cycle, and brings up the magnetic elements that anti-phased with sunspots." + In a recent literature. Thomas and Weiss (2008) proposed a picture of the solar Dynamo on three seales (one large and two small). which. according to the above authors. were only loosely coupled to each other.," In a recent literature, Thomas and Weiss (2008) proposed a picture of the solar Dynamo on three scales (one large and two small), which, according to the above authors, were only loosely coupled to each other." + Ht is not clear if some unknown interplay of different. scale dvnamos may result in the complicated behavior of the Sun's small-scale fields., It is not clear if some unknown interplay of different scale dynamos may result in the complicated behavior of the Sun's small-scale fields. + If we adopt ihe common vision that the smaller magnetic elements are created bv a local turbulent dvnamo. then the local turbulent dvnamo on a certain scale must have closely correlated to the elobal mean-fiekd dynamo.," If we adopt the common vision that the smaller magnetic elements are created by a local turbulent dynamo, then the local turbulent dynamo on a certain scale must have closely correlated to the global mean-field dynamo." + The global dvnaimo either provides seed Παν or moclilies the condition for this turbulent dvnamo., The global dynamo either provides seed flux or modifies the condition for this turbulent dynamo. + At the smallest end. the dynamo is likely to be more'llocal'.," At the smallest end, the dynamo is likely to be more." +. The turbulent. dvnamo. either global or purely local. brings a tremendous amount of turbulent flux to the Sun that continuously interacts with the productis of the mean-field diamo.," The turbulent dynamo, either global or purely local, brings a tremendous amount of turbulent flux to the Sun that continuously interacts with the products of the mean-field dynamo." + The interaction seems to not only help with the operation of the global denamo. but also power the ceaseless siall-scale magnetic activity. aud maintain the Sun'ss Povnting flux to Earth ancl interplanetary space.," The interaction seems to not only help with the operation of the global dynamo, but also power the ceaseless small-scale magnetic activity and maintain the s Poynting flux to Earth and interplanetary space." +"flows the eas is near LOK and the cooling timescale is of order the recombination timescale, zi10fn vr. where s is the hydrogen deusitv (cu?)","flows the gas is near $10^4$ K and the cooling timescale is of order the recombination timescale, $\tau_{\rm rec} \sim 10^5/n$ yr, where $n$ is the hydrogen density $^{-3}$ )." + Since he free-fall timescale is of order 105 vr. the gas loses wdrostatic support aud frec-falls wherever the density exceedsLOn 7," Since the free-fall timescale is of order $10^8$ yr, the gas loses hydrostatic support and free-falls wherever the density $\sim 10^{-3}$ $^{-3}$." +" In the halos we are to consider, his occurs at about 6 scale radii or 50 kpc."," In the halos we are to consider, this occurs at about 6 scale radii or 50 kpc." +" Within his radius, the eas is effectively pressure-free and in yee-fall. achieving velocities of several hundred kn s1 w the time it reaches the halo core."," Within this radius, the gas is effectively pressure-free and in free-fall, achieving velocities of several hundred km $^{-1}$ by the time it reaches the halo core." +" Iu this picture of rapid accretion onto a dense ouwvonie core im the dark matter halo. the production of ionizing photons, Vy is given by: where M ds the mass accretion rate. ipods the mean atomic weight of the infalling gas. gp is the mass of a hydrogen atom."," In this picture of rapid accretion onto a dense baryonic core in the dark matter halo, the production of ionizing photons, $\dot{N}$, is given by: where $\dot{\cal M}$ is the mass accretion rate, $\mu$ is the mean atomic weight of the infalling gas, $m_{\rm H}$ is the mass of a hydrogen atom." +" For cach halo, we therefore need to compute both AM and the velocity of accretion onto the barvouic core, Which by use of equation 1. will then give us οην."," For each halo, we therefore need to compute both $\dot{\cal M}$ and the velocity of accretion onto the baryonic core, which by use of equation \ref{phi} will then give us $ \phi(v)_{\rm up}$." +" For the halos. we use the Einusato density profile (Navarroetal.200:Springel2008) for tle DAL density. in ternis of a scaled radius «6—rfr Iu this equation. the logarithinic derivative is taken tobeequalto 2 at the scale radius, so that p(1)=p»."," For the halos, we use the Einsato density profile \citep{Navarro04,Springel08} for the DM density, in terms of a scaled radius $x = r/r_s$: In this equation, the logarithmic derivative is taken to be equal to $-2$ at the scale radius, so that $\rho(1) = \rho_{-2}$." +" The parameter a has a sinall variation about à=0.15. which is the value adopted here; and p2 is chosen to eive the halo mean density insice the virial radius, pu. once the physical virial radius and halo overdeusitv are known."," The parameter $\alpha$ has a small variation about $\alpha = 0.18$, which is the value adopted here, and $\rho_{-2}$ is chosen to give the halo mean density inside the virial radius, $\rho_{\rm halo}$, once the physical virial radius and halo over–density are known." +" The overall scaling is obtained by requiring pua to be related to the critical density of the Universe at the redshift considered, Assuniug the collapse ofa spherical top hat perturbation the ratio of the mean halo deusity to the critical density of the Universe, A.=Prats/p.. is given by Norman LOO8):: where We take Qa,=0.2715 and O4=0.728 from the fy maxima likelihood value in the WALIAP seven vear analysis (I&omiatsuetal2011)."," The overall scaling is obtained by requiring $\rho_{\rm halo}$ to be related to the critical density of the Universe at the redshift considered, Assuming the collapse of a spherical top hat perturbation, the ratio of the mean halo density to the critical density of the Universe, $\Delta_c = \rho_{\rm halo}/\rho_{c}$, is given by \citep{Bryan98}; where We take $ \Omega_M = 0.2715$ and $ \Omega_\Lambda = 0.728$ from the $H_0$ maximum likelihood value in the WMAP seven year analysis \citep{Komatsu11}." +". The value of the virial radius ayy. iu units of the scale radius is obtained usine a fit bv Dutyetal.(2005). (termed as concentration in that worl) with the desired virial mass for the appropriate redshift +, and cosinoloey (here /—0.701) 1 Once the virial radius is known in terms of the scale radius, the plysical scale is chosen so that the mass of the halo averaged over the virial volume gives the desired mean halo deusitv."," The value of the virial radius $x_{\rm vir}$ in units of the scale radius is obtained using a fit by \citet{Duffy08}, (termed as concentration in that work) with the desired virial mass for the appropriate redshift $z$, and cosmology (here $h = 0.704$ ): Once the virial radius is known in terms of the scale radius, the physical scale is chosen so that the mass of the halo averaged over the virial volume gives the desired mean halo density." + The gravitational potential is then calculated integrating (7)=GAI(r)/r., The gravitational potential is then calculated integrating $\Phi(r) = GM(r)/r$. + An analytical form for this potential can be readily derived., An analytical form for this potential can be readily derived. +" As an approximation of the complex plivsics involved in the collapse process, we assunie that the barvonic uatter follows the dark matter until virialization of the dark matter halo, at which point the barvonic matter »eeius to fall freely towards the ceutre of the halo."," As an approximation of the complex physics involved in the collapse process, we assume that the baryonic matter follows the dark matter until virialization of the dark matter halo, at which point the baryonic matter begins to fall freely towards the centre of the halo." +" Thus, in our initial configuration, the barvonic nass starts off roni rest and with the same radial density profile as he dark matter halo."," Thus, in our initial configuration, the baryonic mass starts off from rest and with the same radial density profile as the dark matter halo." +" The barvonic matter is therefore already. bound within the potential of the dark matter malo, and therefore auv kinetic energy of iufall to the initial radius dunrug the assembly of the halo is effective asstuned to have been already dissipated - presumably hroughthe low velocity shocks associated with the halo asseimiblv in the cold accretion scenario."," The baryonic matter is therefore already bound within the potential of the dark matter halo, and therefore any kinetic energy of infall to the initial radius during the assembly of the halo is effective assumed to have been already dissipated - presumably through the low velocity shocks associated with the halo assembly in the cold accretion scenario." +" From this (somewhat uncertain) approximation to he initial state, we have followed the subsequent accretion in 1-D spherical coordinates using the Lagerangiauli warodvuamic code. (Sutherland.2010"," From this (somewhat uncertain) approximation to the initial state, we have followed the subsequent accretion in 1-D spherical coordinates using the Lagrangian hydrodynamic code, \citep{Sutherland10}." +) This code includes all relevant gas plivsics aud eravitaional fields., This code includes all relevant gas physics and gravitational fields. + By using 5 nested comaius. a final resolution of 43 pc is achieved in this simulation.," By using 5 nested domains, a final resolution of 0.3 pc is achieved in this simulation." + This should be compared with the ~0.5 kpe spatial resolution achieved in the best smoothed particle bydrodvuamics (SPIT) cosinological codes., This should be compared with the $\sim 0.5$ kpc spatial resolution achieved in the best smoothed particle hydrodynamics (SPH) cosmological codes. + As expected. within the critical radius at which ff)essure support is lost due to fast cooling (within about Ro~50 kpe as estimated above) the barvonic latter effectively free-fall towards the centre. since it has no thermal pressure support. am very little rotation.," As expected, within the critical radius at which pressure support is lost due to fast cooling (within about $R \sim 50$ kpc as estimated above) the baryonic matter effectively free-falls towards the centre, since it has no thermal pressure support, and very little rotation." +" In the vicinity of the nucleus, the infalline gas is shocked. and collects into a s11all. dense. rotationallv-supported disk whose size is determined by the torquing of the parent halo."," In the vicinity of the nucleus, the infalling gas is shocked, and collects into a small, dense, rotationally-supported disk whose size is determined by the torquing of the parent halo." + The vertical scale height of this disk is determined by thermal and turbulent support., The vertical scale height of this disk is determined by thermal and turbulent support. +" Its pressure ids given by the ram pressure of the accretion flow, since the radiative shocks can be treated"," Its pressure is given by the ram pressure of the accretion flow, since the radiative shocks can be treated" +twelve dises.,twelve discs. + He finds ωςρω~ GBELO per cent. below maximal according to the definition of Sackett (1997).," He finds $v_{stars}/v_{total} \sim$ $\pm$ 10 per cent, below maximal according to the definition of Sackett (1997)." + The small variation in. disc contribution to. the 2237|0305 rotation curve. as demonstrated in Figure 2.. provides a &ood determination of the degree of maximality in this galaxy.," The small variation in disc contribution to the 2237+0305 rotation curve, as demonstrated in Figure \ref{rot_curve}, provides a good determination of the degree of maximality in this galaxy." + The overall rotation curve. although not constrained: observationally in the region. where this calculation is made (the disc maximum at à 2.2r4). is also reasonably tight given the HE constraints and the profiles we have used.," The overall rotation curve, although not constrained observationally in the region where this calculation is made (the disc maximum at $r$ $\sim$ $r_d$ ), is also reasonably tight given the HI constraints and the profiles we have used." + ‘The contribution of the disc to the rotation has already been determined. by its mean central surface mass density (de = 506430 M. 7j., The contribution of the disc to the rotation has already been determined by its mean central surface mass density (dc = $\pm$ 30 $_\odot$ $^{-2}$ ). +" ""This corresponds to a maximum rotation of Paps,(2. = 168+5 ", This corresponds to a maximum rotation of $v_{disc}(2.2r_d)$ = $\pm$ 5 $^{-1}$. +The maximum rotation is calculated to be (iint: = 2ss+5 , The maximum rotation is calculated to be $v_{total}(2.2r_d)$ = $\pm$ 5 $^{-1}$. +"The percentage contribution of the dise to the rotation. the degree. is therefore. ""his value fits well with that found by Bottema (1993). and is well defined for the potential solutions presented in this work."," The percentage contribution of the disc to the rotation, the degree, is therefore, This value fits well with that found by Bottema (1993), and is well defined for the potential solutions presented in this work." + The disc is clearly sub-maximal., The disc is clearly sub-maximal. + The ιν observed from cach image in a gravitationallv lensecL system. is a direct measure of the magnification in that region of the lens plane.," The flux observed from each image in a gravitationally lensed system, is a direct measure of the magnification in that region of the lens plane." + A comparison between the ας ratios of the images observed in 2237|0305 ancl those predicted by the solutions can further act as a check on the results., A comparison between the flux ratios of the images observed in 2237+0305 and those predicted by the solutions can further act as a check on the results. + Observed fluxes in individual images are a combination of magnification due to microlensing. macrolensing ancl intrinsic. variability. coupled: with time delavs.," Observed fluxes in individual images are a combination of magnification due to microlensing, macrolensing and intrinsic variability coupled with time delays." + Agol. Jones Blaes (2000). however. have measured LB. tluxes (5.9 11.) for the four components and from these the ratio of Ηχος can be calculated.," Agol, Jones Blaes (2000), however, have measured IR fluxes (8.9 $\mu$ m) for the four components and from these the ratio of fluxes can be calculated." + In this region of the spectrum. nmücrolensing events are not observed and it is therefore »ostulated. that these observations sample an extended region of the source.," In this region of the spectrum, microlensing events are not observed and it is therefore postulated that these observations sample an extended region of the source." + In addition. the infrared fluxes are not sensitive to the dust reddening cllects of optical light ravelling through the galaxy.," In addition, the infrared fluxes are not sensitive to the dust reddening effects of optical light travelling through the galaxy." + Thus the LR fluxes should measure the macro-maegnification., Thus the IR fluxes should measure the macro-magnification. + These Luxes. relative to he DB image. are displaved in Table 7..," These fluxes, relative to the B image, are displayed in Table \ref{flux}." + The [lux ratios were calculated. from our results. by aking the ratio of the magnification for cach of the images relative to the D image., The flux ratios were calculated from our results by taking the ratio of the magnification for each of the images relative to the B image. + The magnification is calculated. by aking the ratio of areas of triangles around the images mapped. from the image to the source. plane., The magnification is calculated by taking the ratio of areas of triangles around the images mapped from the image to the source plane. + ον were calculated from the best-fit solution of Table 2 ancl are clisplaved in Table 7.., They were calculated from the best-fit solution of Table \ref{param_dev} and are displayed in Table \ref{flux}. + Our results are consistent with the observations., Our results are consistent with the observations. + The halo model used is a generic profile that is analytically. simple but has little physical motivation., The halo model used is a generic profile that is analytically simple but has little physical motivation. + It does. however provide a mass profile with varving slope. and this is a useful attribute if one wishes to study the gradient. of the mass distribution.," It does, however provide a mass profile with varying slope, and this is a useful attribute if one wishes to study the gradient of the mass distribution." + The best-fit solution for the mass distribution of this galaxy is not an adequate fit., The best-fit solution for the mass distribution of this galaxy is not an adequate fit. + The rotation curve is rising ab the outer regions studied: instead. of falling. and the images are not close enough to their measured. points.," The rotation curve is rising at the outer regions studied instead of falling, and the images are not close enough to their measured points." + A halo with a smaller core would alleviate the former problem but require a change of the other mass distributions to address the latter., A halo with a smaller core would alleviate the former problem but require a change of the other mass distributions to address the latter. + The miass-to-light profile of the. disc suggests an increase in the cise scale length., The mass-to-light profile of the disc suggests an increase in the disc scale length. + Such a move would change the shear introduced by the disc., Such a move would change the shear introduced by the disc. + In any case. he core radii found as the best solutions in this study jiwe large core radii r. 13-I16kpc.," In any case, the core radii found as the best solutions in this study have large core radii, $r_c \sim$ 13-16kpc." + These values are not consistent with the cuspy central regions of the CDM woliles., These values are not consistent with the cuspy central regions of the CDM profiles. + In. order to reduce the 47 and [find an adequate solution. we firstly require further rotation data.," In order to reduce the $\chi^2$ and find an adequate solution, we firstly require further rotation data." + Phe more volts we have with good accuracy. the more likely we are o be able to eliminate profiles and find the best. solution.," The more points we have with good accuracy, the more likely we are to be able to eliminate profiles and find the best solution." + Until more data is obtained. it is not worth trving to twist 1f halo profile to fit the parameters better.," Until more data is obtained, it is not worth trying to twist the halo profile to fit the parameters better." + I£ the error vars on the HIE rotation points were doubled. the fit would come quite consistent with the observations.," If the error bars on the HI rotation points were doubled, the fit would become quite consistent with the observations." + The Barnes et al. (, The Barnes et al. ( +1999) observations were undertaken with the VLA C array. a compact configuration with low angular resolution.,"1999) observations were undertaken with the VLA C array, a compact configuration with low angular resolution." + More informative data is attainable with higher resolution observations., More informative data is attainable with higher resolution observations. + The surface mass distribution of the four components combined can be represented on a log-log plot to study the slope as a function of radius., The surface mass distribution of the four components combined can be represented on a log-log plot to study the slope as a function of radius. + This information is clisplaved in Figure 6.., This information is displayed in Figure \ref{slope}. + The dashed lines are. normalised fits to Xxor? (short-dashed) and Xxr (long-dashed)., The dashed lines are normalised fits to $\Sigma \propto r^{-0.3}$ (short-dashed) and $\Sigma \propto r^{-1}$ (long-dashed). + The central slope overall is reasonably steep given the [aree inlluence of the bulge., The central slope overall is reasonably steep given the large influence of the bulge. + The dark matter halo has zero slope in the inner regions given its large core radius., The dark matter halo has zero slope in the inner regions given its large core radius. + The transition to isothermal is expected. given the dominance of the halo in the outer regions., The transition to isothermal is expected given the dominance of the halo in the outer regions. + We have undertaken a study of the structure of galaxy 2237)0305 using cynamical anc gravitational lensing constraints., We have undertaken a study of the structure of galaxy 2237+0305 using dynamical and gravitational lensing constraints. + Phe combination of these techniques allows, The combination of these techniques allows +One wav in which barred spiral galaxies are commonly divided into subclasses is by takine note of where (he large scale spiral starts wilh relation to the bar 1994).,"One way in which barred spiral galaxies are commonly divided into subclasses is by taking note of where the large scale spiral starts with relation to the bar \citep[e.g.,][]{sandage94}." +. In 5D(s) galaxies. the spiral arms begin at the ends of the bar. whereas in SB) galaxies. (he spiral armis beein on a ring connecting the ends of the bar.," In SB(s) galaxies, the spiral arms begin at the ends of the bar, whereas in SB(r) galaxies, the spiral arms begin on a ring connecting the ends of the bar." + 10(19) ealaxies are a transition group., SB(rs) galaxies are a transition group. + SD(s) structures are (thought (to be preferentially found in less strongly barred. galaxies than their 5D(r) counterparts (e.g..Sanders&Tubbs1980:Simkinetal. 1930).," SB(s) structures are thought to be preferentially found in less strongly barred galaxies than their SB(r) counterparts \citep[e.g.,][]{sanders80, simkin80}." +. Furthermore. SD(s) galaxies typically show large-scale dust lanes that are nol present in 5D(r) galaxies. ancl SD(r) galaxies are observed (o have less dust in (heir central regions (han SB(s) galaxies (Ixormendy.&Ixennieutt.2004).," Furthermore, SB(s) galaxies typically show large-scale dust lanes that are not present in SB(r) galaxies, and SB(r) galaxies are observed to have less dust in their central regions than SB(s) galaxies \citep{kormendy04}." +.. This is an apparent inconsistenev with the generic bar-[ueling picture: 5D(r) galaxies are thought to have less dustbut be more strongly barreclwhile more strongly barred galaxies should have a higher central dust content., This is an apparent inconsistency with the generic bar-fueling picture: SB(r) galaxies are thought to have less dust—but be more strongly barred—while more strongly barred galaxies should have a higher central dust content. + As these conclusions have been based on measuring the barstrength as an axis ralio (e.g..Sanders&Tubbs1980).. rather than as a force ratio. this discrepancy might be because SD(r) galaxies are actually more weakly barred than SD(s) galaxies.," As these conclusions have been based on measuring the barstrength as an axis ratio \citep[e.g.,][]{sanders80}, rather than as a force ratio, this discrepancy might be because SB(r) galaxies are actually more weakly barred than SB(s) galaxies." + We reinvestigate the relation between these bar sub-tvpes and barstrength., We reinvestigate the relation between these bar sub-types and barstrength. + The breakdown by nuclear class. including LGD. of the 21 RC3-classified SB galaxies in our sample is given in Table 2..," The breakdown by nuclear class, including LGD, of the 21 RC3-classified SB galaxies in our sample is given in Table \ref{tbl:SB}." + A small increase in ιν is seen from SBir) to 9105)., A small increase in $Q_b$ is seen from SB(r) to SB(s). + Visually. the SD(r) galaxies have much less dust structure than the SD(rs) and SB(s) galaxies.," Visually, the SB(r) galaxies have much less dust structure than the SB(rs) and SB(s) galaxies." + According to the Wilcoxon test. the SD(r) sample has a smaller Qy than the SB(s) galaxies ab a confidence level of," According to the Wilcoxon test, the SB(r) sample has a smaller $Q_b$ than the SB(s) galaxies at a confidence level of." +94%... Sanders&Tubbs(1980). suggest bar pattern speed as an alternative origin to the differences in SD(r) and SD(s) structure: a slowly rotating bar should give rise to 5D(r) structure. while a rapidly rotating bar should vield SB(s) structure.," \citet{sanders80} suggest bar pattern speed as an alternative origin to the differences in SB(r) and SB(s) structure: a slowly rotating bar should give rise to SB(r) structure, while a rapidly rotating bar should yield SB(s) structure." + As the differences in barstrength are reversed from what was expected. rotation marx be more important in determining the large-scale morphology.," As the differences in barstrength are reversed from what was expected, rotation may be more important in determining the large-scale morphology." +" With respect to the amount of dust structure. we find with >99%. confidence that the SBfr) galaxies have less dust structure (smaller 0,4) than the SB(s) sample: similarly. we find with confidence that the SD(r) saniple has less dust than the 9D(1s) galaxies."," With respect to the amount of dust structure, we find with $\ge 99$ confidence that the SB(r) galaxies have less dust structure (smaller $\sigma_{\mbox{\scriptsize sm}}$ ) than the SB(s) sample; similarly, we find with confidence that the SB(r) sample has less dust than the SB(rs) galaxies." +We vow present the proof of Theorem 1.2 that involve finer estimates on the Hólkler For the sake of simplifying the ideas of the proof. we only couskler 1- dimensious «=;04.,"We now present the proof of Theorem \ref{theo2} that involve finer estimates on the Höllder For the sake of simplifying the ideas of the proof, we only consider $1$ -spatial dimensions $x=x_{1}$." + The general n-dimensional case cau be easily deduced., The general $n$ -dimensional case can be easily deduced. +" Following the samenotations of [5].. we let OQ=(—1.2)x(CT.2T). 2,€Z3COy such that and We also take the cut-olf function W€CO(E?) 0€Wx1 satis[ving: The main idea of the proof consists in exteudiug the Muection / to a suitable function of the form Wf wheref is clelined on £24."," Following the samenotations of \cite{Ibrahim09}, we let $\widetilde{\O}_{T} = (-1,2) \times (-T,2T)$, $\mathcal{Z}_{1} +\subseteq \mathcal{Z}_{2} \subseteq \widetilde{\O}_{T}$ such that and We also take the cut-off function $\Psi\in C^{\infty}_{0}(\R^{2})$, $0\leq\Psi\leq 1$ satisfying: The main idea of the proof consists in extending the function $f$ to a suitable function of the form $ \Psi\tilde{f}$ where$\tilde{f}$ is defined on $\widetilde{\O}_{T}$." + We then apply inequality (1.1)) (the scalar-valued. version. with ΗΞ 1) ιο Wf aud we estimate the different norms iu order to get the result.," We then apply inequality \ref{Ib:eq4}) ) (the scalar-valued version with $n=1$ ) to $ +\Psi\tilde{f}$ and we estimate the different norms in order to get the result." + However. away from the complicated extension (Sobolev extension) of the function f that was done in [5].. we here consider a sliupler sviunetric exteusion.," However, away from the complicated extension (Sobolev extension) of the function $\tilde{f}$ that was done in \cite{Ibrahim09}, we here consider a simpler symmetric extension." + Lucleecd. we first take the spatial syiumetry of the function f: aud then the sviminetry. with respect to /: claim that Wf €C7(8?) withΠλ In this case. we apply the scalar-valued version of inequality (1.1)) (see Remark 1.3)) to the function Wf with ;—1 and gGr./)=INWy.fly.Ody.," Indeed, we first take the spatial symmetry of the function $f$: and then the symmetry with respect to $t$ : We claim that $\Psi \tilde{f} \in C^{\g,\g/2}(\R^{2})$ with In this case, we apply the scalar-valued version of inequality \ref{Ib:eq4}) ) (see Remark \ref{rem1}) ) to the function $\Psi +\tilde{f}$ with $i=1$ and $g(x,t) = +\int_{0}^{x} \Psi(y,t) \tilde{f}(y,t) dy$." + This. together with the fact that Y=1 on ο). lead to the following estimate: ]t is worth noticing that choosing 7=1 above is somehow restrictive.," This, together with the fact that $\Psi = 1$ on $\O_T$, lead to the following estimate: It is worth noticing that choosing $i=1$ above is somehow restrictive." + Iu fact. we could also have used the inequality with /—2 and gGr./)—fiWGr.s)fGr. s)ds.," In fact, we could also have used the inequality with $i=2$ and $g(x,t) = \int_{0}^{t} \Psi(x,s) \tilde{f}(x,s) ds$ ." +" In [7]. it was shown that |W||paoce,€CULLBALO(Oy}+ μμ... while it is clear that ΠρCWlexοιSOWPee onu "," In \cite{IM09} it was shown that $\|\Psi \tilde{f}\|_{BMO(\R^{2})} +\leq C (\|f\|_{BMO(\O_{T})} + \|f\|_{L^{1}(\O_{T})})$ , while it is clear that $\|g\|_{L^{\infty}(\R^{2})} \leq C +\|\tilde{f}\|_{L^{\infty}(\widetilde{\O}_T)}\leq C +\|f\|_{C^{\g,\g/2}(\O_{T})}$ ." +These arguments. along with (3.8))and (3.9)). directly terminate the proof.," These arguments, along with \ref{h}) )and \ref{estimate1}) ), directly terminate the proof." + The ouly point left is to show the claim (3.8))., The only point left is to show the claim\ref{h}) ). + Recall the norm, Recall the norm +which is not surprising given the increased resolution.,which is not surprising given the increased resolution. +" Observations indicate that the binary period distribution is extremely broad, covering separations from only a few stellar radii to >104 AU (?).."," Observations indicate that the binary period distribution is extremely broad, covering separations from only a few stellar radii to $\ga 10^4$ AU \citep{duquennoy91a}." +" It is therefore not surprising that some binaries that might be resolved into two separate stars in run HR instead appear as a single star in run LR - indeed, we would expect this result in essentially any simulation that did not resolve the radii of individual stars."," It is therefore not surprising that some binaries that might be resolved into two separate stars in run HR instead appear as a single star in run LR – indeed, we would expect this result in essentially any simulation that did not resolve the radii of individual stars." +" Nonetheless, notice that, if we normalize to the number of stars present at equal times and fractions of mass accreted, then the difference between the two runs disappears."," Nonetheless, notice that, if we normalize to the number of stars present at equal times and fractions of mass accreted, then the difference between the two runs disappears." + The number of stars present at any given time in run HR is roughly 1.6 times the number present at the same time in run LR., The number of stars present at any given time in run HR is roughly $1.6$ times the number present at the same time in run LR. +" Thus the trend in terms of when the stars are formed in the simulations is nearly identical in the two cases, and we can regard as well-resolved the distribution in time of when stars form."," Thus the trend in terms of when the stars are formed in the simulations is nearly identical in the two cases, and we can regard as well-resolved the distribution in time of when stars form." + The trend of number of stars versus mass shown in Figure 5 is interesting., The trend of number of stars versus mass shown in Figure \ref{starhist2} is interesting. +" In the radiative runs, when M.ox/ΜεS;0.1, the number of stars increases roughly linearly with the total stellar mass, as we might expect if the mass per star were constant."," In the radiative runs, when $M_{*,\rm tot}/M_c \la 0.1$, the number of stars increases roughly linearly with the total stellar mass, as we might expect if the mass per star were constant." +" However, the rate at which new stars appears drops sharply once M, 0.2."," However, the rate at which new stars appears drops sharply once $M_{*,\rm tot}/M_c \ga 0.2$ ." +" Indeed, we see that 60—7096 of all stars have totformed/Me at a time when only ~10% of the cloud mass has been incorporated into stars, By the time of the cloud mass has gone into stars, nearly of all the stars are in place."," Indeed, we see that $60-70\%$ of all stars have formed at a time when only $\sim 10\%$ of the cloud mass has been incorporated into stars, By the time of the cloud mass has gone into stars, nearly of all the stars are in place." +" In effect, the fragmentation of the gas into new stars has completely shut down."," In effect, the fragmentation of the gas into new stars has completely shut down." +" Given that this effect occurs nearly identically in runs LR and HR, this cannot be a resolution effect."," Given that this effect occurs nearly identically in runs LR and HR, this cannot be a resolution effect." +" In contrast, run ISO shows very different behavior."," In contrast, run ISO shows very different behavior." +" The number of stars as a function of total stellar mass is almost the same as in run HR up to the point where ~15% of the mass has been incorporated into stars, but the two runs diverge after that."," The number of stars as a function of total stellar mass is almost the same as in run HR up to the point where $\sim 15\%$ of the mass has been incorporated into stars, but the two runs diverge after that." +" New stars continue forming all the way through run ISO, at a rate that is only slightly less after M,tor/M_.20.2 than it was earlier in the simulation."," New stars continue forming all the way through run ISO, at a rate that is only slightly less after $M_{*,\rm tot}/M_c \ga 0.2$ than it was earlier in the simulation." +" This strongly suggests that the shutdown in new star formation we observe in runs LR and HR is a radiative effect, a topic to which we will return in Section 3.3.."," This strongly suggests that the shutdown in new star formation we observe in runs LR and HR is a radiative effect, a topic to which we will return in Section \ref{sec:thermo}." +" As one might expect, this shutoff of fragmentation into new stars in runs LR and HR even as the total stellar mass continues to increase produces a dramatic effect on the stellar mass distribution."," As one might expect, this shutoff of fragmentation into new stars in runs LR and HR even as the total stellar mass continues to increase produces a dramatic effect on the stellar mass distribution." + Figures 6 and 7 show the cumulative and differential mass distributions of the stars formed in our simulations at the times when the total mass in stars is 10—5096 of the initial cluster mass., Figures \ref{imfplot1} and \ref{imfplot2} show the cumulative and differential mass distributions of the stars formed in our simulations at the times when the total mass in stars is $10-50\%$ of the initial cluster mass. + All these plots show that the stellar mass distribution in the radiative runs moves continuously to higher masses as the simulation proceeds., All these plots show that the stellar mass distribution in the radiative runs moves continuously to higher masses as the simulation proceeds. +" This is because mass is accreting onto existing stars, which rise in mass, but very few new, lower-mass stars are forming."," This is because mass is accreting onto existing stars, which rise in mass, but very few new, lower-mass stars are forming." +" Note that, while the mean stellar masses are slightly different in runs LR and HR, the systematic drift of these mean to higher masses as the total stellar mass rises appears to about occur equally in"," Note that, while the mean stellar masses are slightly different in runs LR and HR, the systematic drift of these mean to higher masses as the total stellar mass rises appears to about occur equally in" +"The best-fit solutions of the Monte Carlo samples (dots in panels c to h, Fig. 9))","The best-fit solutions of the Monte Carlo samples (dots in panels c to h, Fig. \ref{fig:fit_sn5}) )" +" are of course centered around the true solution, which explains the small relative shift of the x-contours, which indicate the best-fit solutions for the simulated data."," are of course centered around the true solution, which explains the small relative shift of the $\chi^2$ -contours, which indicate the best-fit solutions for the simulated data." + As statistically expected this shift is smaller than the lo error., As statistically expected this shift is smaller than the $\sigma$ error. + In the second test case (Fig. 10)), In the second test case (Fig. \ref{fig:fit_sn1.5}) ) +" we considered only 50, hhalf as many, data points and we considerably reduced the signal-to-noise ratio down to 1.5, which is more representative of the data by ?.."," we considered only 50, half as many, data points and we considerably reduced the signal-to-noise ratio down to 1.5, which is more representative of the data by \citet{berdyuginaetal2008a}." + Here we assumed a circular orbit so that the longitude of the periastron becomes undefined., Here we assumed a circular orbit so that the longitude of the periastron becomes undefined. +" For the x?-minimization we nonetheless kept the eccentricity e as a free parameter, but we fixed w to90°.."," For the $\chi^2$ -minimization we nonetheless kept the eccentricity $e$ as a free parameter, but we fixed $\omega$ to." + For very small eccentricities ω is not well constrained by the data but has also negligible influence on the polarization curves., For very small eccentricities $\omega$ is not well constrained by the data but has also negligible influence on the polarization curves. +" Therefore, we are left with six free parameters."," Therefore, we are left with six free parameters." + Even under these more difficult circumstances the fitting procedure proved to be very robust., Even under these more difficult circumstances the fitting procedure proved to be very robust. + The original curves are again quite well reproduced (Fig., The original curves are again quite well reproduced (Fig. + 10aa&bb) and the input parameters well identified (Table 1))., \ref{fig:fit_sn1.5}a b) and the input parameters well identified (Table \ref{table:parameters}) ). + The y?-minimum, The $\chi^2$ -minimum +the cross power spectra from the overlapping portion of resicuals of cach pulsar pair with no further. processing.,the cross power spectra from the overlapping portion of residuals of each pulsar pair with no further processing. + Llowever. upon simulating this procedure. we found that the lowest frequencies in the cross power spectra were biasec whenever diiJdosgsnp.," However, upon simulating this procedure, we found that the lowest frequencies in the cross power spectra were biased whenever $\tobs > \tol$." + This bias took the form of a significantly. non-zero imaginary part in the cross power spectrum., This bias took the form of a significantly non-zero imaginary part in the cross power spectrum. + Also. we found that much of the correlated signa at low frequencies was removed. as shown in Figure 7..," Also, we found that much of the correlated signal at low frequencies was removed, as shown in Figure \ref{fig:fit}." + We were unable to eliminate these cllects unless we performer a WLSO fit of a quadratic function for each time series over the overlapping time range., We were unable to eliminate these effects unless we performed a WLSQ fit of a quadratic function for each time series over the overlapping time range. + Γι restores the correlation in the GWD signal between dillerent pulsars (right panels of Figure 7))., This restores the correlation in the GWB signal between different pulsars (right panels of Figure \ref{fig:fit}) ). + This additional WLSQ fit will introduce a new bias because of removing some of the GWD signal at fMua. but this new bias is easily corrected with the calibration factors 5;;Cf£).," This additional WLSQ fit will introduce a new bias because of removing some of the GWB signal at $f=1/\tol$, but this new bias is easily corrected with the calibration factors $\gamma_{ij}(f)$." + However. there is an additional loss of 10 per cent of the GWD signal in the Verbiest ct al. (," However, there is an additional loss of 10 per cent of the GWB signal in the Verbiest et al. (" +2005. 2009) observations because ofthis extra WLSQ fit.,"2008, 2009) observations because of this extra WLSQ fit." + The CAVB analysis is complicated by the unknown cllects of other correlated signals in the timing residuals., The GWB analysis is complicated by the unknown effects of other correlated signals in the timing residuals. + Instabilities in TPP and errors in the Solar-Systemi ephemeris both produce signals which are correlated between: dillerent pulsars., Instabilities in TT and errors in the Solar-System ephemeris both produce signals which are correlated between different pulsars. + We estimated the effect. of these uncertainties by using an updated: timescale anc the most recent Solar-System ephemeris., We estimated the effect of these uncertainties by using an updated timescale and the most recent Solar-System ephemeris. + Instabilities in PP produce a positive cross correlation independent of angular separation., Instabilities in TT produce a positive cross correlation independent of angular separation. + Any estimate of the clock error will thus be correlated with the estimate of the GWB amplitude., Any estimate of the clock error will thus be correlated with the estimate of the GWB amplitude. + Lad. we made a significant detection of the GWD. this would. have to be accounted for.," Had we made a significant detection of the GWB, this would have to be accounted for." + ‘To estimate the importance of possible clock instabilities. we processed the Verbiest et al. (," To estimate the importance of possible clock instabilities, we processed the Verbiest et al. (" +2008. 2009) observations using the version of TE. released by DIPM. in 2010 (sec.e.g. 7)..,"2008, 2009) observations using the version of TT released by BIPM in 2010 \citep[see, e.g.,][]{2003Petit}." + his post-corrected timescale has revealed. statistically in T, This post-corrected timescale has revealed statistically significant inaccuracies in TT(TAI). +able 6.., The results are shown in Table \ref{tbl:ttephem}. + While the change of clock reference only changes our estimated GAB level by nine per cent of the uncertainty. the absolute change (0.810 77) is at a significant level for some predictions of the GAVB (??)..," While the change of clock reference only changes our estimated GWB level by nine per cent of the uncertainty, the absolute change $0.8\e{-30}$ ) is at a significant level for some predictions of the GWB \citep{jb03,svc08}." + This implies that such instabilities in LE must be accounted for when analysing future data sets., This implies that such instabilities in TT must be accounted for when analysing future data sets. + The results. from using. the newest Solar-Svsteni ephemeris DE421 (2). are given in Table 6.., The results from using the newest Solar-System ephemeris DE421 \citep{2009DE421} are given in Table \ref{tbl:ttephem}. + While there have been some improvements in this ephemeris version compared to DEL05. most of the changes are absorbed by the pulsar parameter fit.," While there have been some improvements in this ephemeris version compared to DE405, most of the changes are absorbed by the pulsar parameter fit." + The estimated GWD level has changed by 24 per cent of the uncertainty., The estimated GWB level has changed by 24 per cent of the uncertainty. + Lowe assume DIZ421 is correct. then the use of DIZ405 is similar to introducing a spurious CAB signal with 44—L5.10C. a signal which is undetectable in most time series from the Verbiest ct al. (," If we assume DE421 is correct, then the use of DE405 is similar to introducing a spurious GWB signal with $A=1.5\e{-15}$, a signal which is undetectable in most time series from the Verbiest et al. (" +2008. 2009) observations.,"2008, 2009) observations." + However. future observations will need to account [or the cllects of inaccuracies in the Solar-System ephemoris.," However, future observations will need to account for the effects of inaccuracies in the Solar-System ephemeris." + ]t is dillicult to. determine. the exact. contributions to the weighting of cach pulsar pair when using error. bars derived from Monte Carlo simulations., It is difficult to determine the exact contributions to the weighting of each pulsar pair when using error bars derived from Monte Carlo simulations. + Phe dominant cllect is the size of Docertap., The dominant effect is the size of $\tol$. + For a GWB caused. by SMDILDDs. the weighting factor increases approximately as TAS.," For a GWB caused by SMBHBs, the weighting factor increases approximately as $\tol^{4.3}$." + A higher noise level in the residuals of cach pulsar in the pair will decrease the weight of that pair approximately linearly., A higher noise level in the residuals of each pulsar in the pair will decrease the weight of that pair approximately linearly. +" The angle subtended at the observer by the pair of pulsars 4), can be important 1£8;; is near the zeroes ofthe function plotted in Figure 1..", The angle subtended at the observer by the pair of pulsars $\theta_{ij}$ can be important if $\theta_{ij}$ is near the zeroes of the function plotted in Figure \ref{fig:HD}. . + ‘To determine which pulsars contribute the most to our estimate of the GCWD. we perform the WLSQ fit described by Equations (10)) and (11)) to only 189 of the possible 190 (6;;) estimates.," To determine which pulsars contribute the most to our estimate of the GWB, we perform the WLSQ fit described by Equations \ref{eq:a2est}) ) and \ref{eq:a2esterr}) ) to only 189 of the possible 190 $\asqzeta$ estimates." + By varving which estimate of Anc(6;;) is Anremoved. we can find the pulsar pairs which have the ereatest influence over the measurement. of 217. in these residuals.," By varying which estimate of $\asqzeta$ is removed, we can find the pulsar pairs which have the greatest influence over the measurement of $\hat{A^2}$ in these residuals." + This is performed by finding Ald? for cach pair of pulsars. whichis the measured? from all pulsar pairs minus the value of AP when including the given pulsar pair.," This is performed by finding $\Delta\hat{A^2}$ for each pair of pulsars, which is the measured $\hat{A^2}$ from all pulsar pairs minus the value of $\hat{A^2}$ when including the given pulsar pair." + Those pairs with the lareest contribution to this measure are given in Table 7.. and a histogram of the absolute value AL for all pulsar pairs is provided in Figure &..," Those pairs with the largest contribution to this measure are given in Table \ref{tbl:psrpairs}, and a histogram of the absolute value $\left|\Delta\hat{A^2}\right|$ for all pulsar pairs is provided in Figure \ref{fig:histdelA2}." + This analvsis shows that the measurement of 20. is determined by only ai few pulsar pairs., This analysis shows that the measurement of $\hat{A^2}$ is determined by only a few pulsar pairs. + This. severely reduces the number of degrees of freedom. when detecting the GAB. and thus decreases the maximum. attainable detection confidence (see?) because it reduces our ability to average out the sell-noise in the residuals caused by the CAB signal at each. pulsar.," This severely reduces the number of degrees of freedom when detecting the GWB, and thus decreases the maximum attainable detection confidence \citep[see][]{jhlm05} because it reduces our ability to average out the self-noise in the residuals caused by the GWB signal at each pulsar." + Observing morestrong pulsars is essential to increasing the number of degrees of freedom, Observing morestrong pulsars is essential to increasing the number of degrees of freedom +Stars of low and intermediate mass with initial masses between 0.8. —SAL. evolve to the asymptotic giant branch (AGB).,Stars of low and intermediate mass with initial masses between $M_\odot$ – $M_\odot$ evolve to the asymptotic giant branch (AGB). + Then. thanks to severe mass loss. the AGB star evolves rapidly at nearly constant luminosity to higher effective temperatures to the white dwarf cooling track.," Then, thanks to severe mass loss, the AGB star evolves rapidly at nearly constant luminosity to higher effective temperatures to the white dwarf cooling track." + Typical stellar lifetimes of post-AGB stars are expected to be of the order of 10! years (Schónberner 1983)., Typical stellar lifetimes of post-AGB stars are expected to be of the order of $^{\rm 4}$ years nberner 1983). + The gas lost by the AGB star forms a circumstellar shell., The gas lost by the AGB star forms a circumstellar shell. + When ye post-AGB star is cool. the dust in the shell heated by stellar A4o4radiation provides an infrared excess.," When the post-AGB star is cool, the dust in the shell heated by stellar radiation provides an infrared excess." + When the star has traversed the top of the H-R diagram to higher effective temperatures. the circumstellar gas is ionized.," When the star has traversed the top of the H-R diagram to higher effective temperatures, the circumstellar gas is ionized." + Then. the star is said to have evolved to reached the proto-planetary nebula stage.," Then, the star is said to have evolved to reached the proto-planetary nebula stage." + Shortly after this. the post-AGB star has evolved to become a planetary nebula with a hot white dwarf as the central star.," Shortly after this, the post-AGB star has evolved to become a planetary nebula with a hot white dwarf as the central star." + Determinations of the chemical composition for post-AGB star hold the potential of yielding insights into the chemical history of the AGB star and its conversion by mass loss to its slimmer post-AGB form., Determinations of the chemical composition for post-AGB star hold the potential of yielding insights into the chemical history of the AGB star and its conversion by mass loss to its slimmer post-AGB form. + In this paper we present a determination of the chemical composition of IRAS [8095+2704. a post-AGB star with a substantial dusty circumstellar shell.," In this paper, we present a determination of the chemical composition of IRAS 18095+2704, a post-AGB star with a substantial dusty circumstellar shell." + The discovery of the optical counterpart TIRAS 1809542704. was made by Hrivnak. Kwok. Volk (1987. 1988).," The discovery of the optical counterpart IRAS 18095+2704 was made by Hrivnak, Kwok, Volk (1987, 1988)." + This V = 10.4 mag star is a high-latitude F supergiant with a large far-IR excess., This $V$ = 10.4 mag star is a high-latitude F supergiant with a large far-IR excess. + In the assembled by Hrivnak. Kwok. Volk (1988). the star has a peculiar TR continuum slope at wavelengths shortward of the 10 Silicate emission feature.," In the assembled by Hrivnak, Kwok, Volk (1988), the star has a peculiar IR continuum slope at wavelengths shortward of the 10 silicate emission feature." + According to Volk Kwok (1987). this peculiar continuum shape is a result of a detached dust shell.," According to Volk Kwok (1987), this peculiar continuum shape is a result of a detached dust shell." + Observational evidence for an expanding shell came from Lewis. Eder. Terzian (1985) and Eder. Lewis. Terzian (1988) via detection of OH maser emission at [612 and 1665/67 MHz from the Arecibo telescope.," Observational evidence for an expanding shell came from Lewis, Eder, Terzian (1985) and Eder, Lewis, Terzian (1988) via detection of OH maser emission at 1612 and 1665/67 MHz from the Arecibo telescope." + Gledhill et al. (, Gledhill et al. ( +2001) from imaging polarimetry report an extended envelope or a reflection nebula around the star.,2001) from imaging polarimetry report an extended envelope or a reflection nebula around the star. + (0.2 em The pioneering study of TRAS 18095427048 composition was reported by Klochkova (1995) from echelle spectra (2?=24 000)," 0.2 cm The pioneering study of IRAS 18095+2704's composition was reported by Klochkova (1995) from echelle spectra $R += 24\,000$ )" +The Ixuiper Belt is a vast swarm of iev bodies bevond the orbit of Neptune in our solar svstem.,The Kuiper Belt is a vast swarm of icy bodies beyond the orbit of Neptune in our solar system. + Following the discovery of the first INuiper Belt objects (IXDOs) in 1930 (Pluto: and 1992 (1992OD:Jewitt&Luu1993).. several groups began survevs to characterize the limits of the Kuiper Belt," Following the discovery of the first Kuiper Belt objects (KBOs) in 1930 \citep[Pluto;][]{tom46} + and 1992 \citep[1992 QB$_1$;][]{jew93}, several groups began large-scale surveys to characterize the limits of the Kuiper Belt" +5truciu BStrueiu Cnr Cnr Cnr Cnr σα1) clussilü ον CLUSSS ciusshbxlO 1102 ciutir clurd cutis clures 12pt,"5truein 8truein cmr8 cmr8 cmr8 cmr8 cmr10 cmssi10 cmss10 cmss8 cmssbx10 2 cmti7 cmr6 cmti8 cmr8 \def\ref{\par\noindent\hangindent 15pt} + = 12pt" +5truciu BStrueiu Cnr Cnr Cnr Cnr σα1) clussilü ον CLUSSS ciusshbxlO 1102 ciutir clurd cutis clures 12pt-,"5truein 8truein cmr8 cmr8 cmr8 cmr8 cmr10 cmssi10 cmss10 cmss8 cmssbx10 2 cmti7 cmr6 cmti8 cmr8 \def\ref{\par\noindent\hangindent 15pt} + = 12pt" +weights due to having cmereccd roni differeut depths iu the nebula.,weights due to having emerged from different depths in the nebula. + We have ested this method of calculating uncertaimtiCR by running the same model multiple times. but with a different. random uunber seed each time the model was run.," We have tested this method of calculating uncertainties by running the same model multiple times, but with a different random number seed each time the model was run." + This allowed us to compue the true uncertaiutv lu au outpit quantity directly from the variation of the quantity between imodel rus., This allowed us to compute the true uncertainty in an output quantity directly from the variation of the quantity between model runs. + By comparing the truc value of the uncertainty with the values computed usii18o eqs., By comparing the true value of the uncertainty with the values computed using eqs. + 25 and 27.. we were able to evaluate the accuracy of our smnple method of estimating uucertainties m output quautitics.," \ref{eq_unc} and \ref{eq_unc_image}, we were able to evaluate the accuracy of our simple method of estimating uncertainties in output quantities." + For this testing. we chose to use a sphere with a homogeneous dust distribution. a 7:=1l. V baud Milky Wav dust erai properties. and a central ilbhuuinatii18o star (sec Fie.," For this testing, we chose to use a sphere with a homogeneous dust distribution, a $\tau_V = 1$, V band Milky Way dust grain properties, and a central illuminating star (see Fig." + 2aa)., \ref{fig_exam}a a). + Other optical depths eive similar results., Other optical depths give similar results. + We rau the model 100 times and varied the total munhber of photous between 107 and 107., We ran the model 100 times and varied the total number of photons between $10^2$ and $10^5$. + Figure 1 cisplavs the uncertainties in the scattered flux and Q component of the polarized flux as a fiction of the nunber of photous run., Figure \ref{fig_test_unc} displays the uncertainties in the scattered flux and Q component of the polarized flux as a function of the number of photons run. + The general trend is for the muacertaity calculated using eq., The general trend is for the uncertainty calculated using eq. + 25 or eq., \ref{eq_unc} or eq. + 27. to uuderestinate he actual uncertantv by smaller amounts as the uuuer of photons run increases., \ref{eq_unc_image} to underestimate the actual uncertainty by smaller amounts as the number of photons run increases. + This is due to small umuber statistics. especially when the uncertaiutyv was calculatexl using eq. 27..," This is due to small number statistics, especially when the uncertainty was calculated using eq. \ref{eq_unc_image}." + For à large πα: of photous (e.g... 107). he uncertainty calculated using eq.," For a large number of photons (e.g., $10^5$ ), the uncertainty calculated using eq." + 25. or 27 Is a τον eood estimate of the actual uncertainty., \ref{eq_unc} or \ref{eq_unc_image} is a very good estimate of the actual uncertainty. + The reason is tlat he majority of the scattered flax aud Q component of 1C xlarized flux comes from the ceutral region of the nella (sce Fig., The reason is that the majority of the scattered flux and Q component of the polarized flux comes from the central region of the nebula (see Fig. + 2aa) aud the intrinsic variation of the scatter Hux in the ceutral region is small., \ref{fig_exam}a a) and the intrinsic variation of the scattered flux in the central region is small. + This implies that nucertaity calculated frou eq., This implies that the uncertainty calculated from eq. + 25 or 27 is domina o» Monte Carlo noise for this model., \ref{eq_unc} or \ref{eq_unc_image} is dominated by Monte Carlo noise for this model. + There are mo systems where this will not be the case and. as a resu care mnust be taken calculating the uncertainty using uecthod outlined above.," There are model systems where this will not be the case and, as a result, care must be taken calculating the uncertainty using the method outlined above." + We ested the results of the DIRTY model against those produced by Moute Carlo radiative trauster models which do not weight photous aud models which use the Witt(1977) photon weightine., We tested the results of the DIRTY model against those produced by Monte Carlo radiative transfer models which do not weight photons and models which use the \citet{wit77} photon weighting. + These models include ones which we have code as well oues others have coded BBjorkinan 1999. 4mivate commuiuication: WWood 1999. private communication).," These models include ones which we have coded as well ones others have coded Bjorkman 1999, private communication; Wood 1999, private communication)." + For computational reasons. these iiodels are usually resvicted to spherically svuuuetrie systems with smoothly varving radial dust distributions.," For computational reasons, these models are usually restricted to spherically symmetric systems with smoothly varying radial dust distributions." + Our two main test cases were for 7=1 and T== LO.," Our two main test cases were for $\tau = +1$ and $\tau = 10$ ." + We adopted an albedo of 1.6 and a scattering phase fiction asviunetrv of (1.6., We adopted an albedo of 0.6 and a scattering phase function asymmetry of 0.6. + Tn all cases. he YusefZadeh.Morris.&White{1981) photon weighting method produced statistically simula results to the other two weighting methods.," In all cases, the \citet{yus84} + photon weighting method produced statistically similar results to the other two weighting methods." + Tn adeditiou. we computed the wavelength depeudence of the polarization for active ealactic imcleus models simular to those used by. Manzi&diSerego.Alighieri—(199(i) aud found qualitative agreement with their results.," In addition, we computed the wavelength dependence of the polarization for active galactic nucleus models similar to those used by \citet{man96} and found qualitative agreement with their results." + Quantitative agreement is more difficult to test as we used a different «ust era mode than Manzini&diSerego, Quantitative agreement is more difficult to test as we used a different dust grain model than \citet{man96}. +Alighieri)(1996). Figure 2 illustrates the mages produced bx the DIRTY model., Figure \ref{fig_exam} illustrates the images produced by the DIRTY model. + Figure 2aa shows how a spherical ucmula with a central ilhuuimatius star would look iu the V bond assunune Milky Way type ¢lust with a homogeneous distribution and a radial τι=l., Figure \ref{fig_exam}a a shows how a spherical nebula with a central illuminating star would look in the V band assuming Milky Way type dust with a homogeneous distribution and a radial $\tau_V = 1$. + Figure 2bb shows how a biconical nebula or active galactic nucleus inclined by an anele of 30° would look in the V band assuming Milky Way type dust with a homogeneous cüstribution and a τι=1., Figure \ref{fig_exam}b b shows how a biconical nebula or active galactic nucleus inclined by an angle of $\degr$ would look in the V band assuming Milky Way type dust with a homogeneous distribution and a $\tau_V = 1$. + Figure 3. displays the spectra enerev distribution before and after the inclusion of dust i 1a sinple starburst svsten., Figure \ref{fig_exam_sed} displays the spectral energy distribution before and after the inclusion of dust in a simple starburst system. + This is a good illustration of how the dust redistributes energev froii the ultraviolet to the infrared., This is a good illustration of how the dust redistributes energy from the ultraviolet to the infrared. + Iu. addition. the three components (thera equilibrium. thermal nou- aud aromatic feature emissions) of the dust Cluission spectrum are shown.," In addition, the three components (thermal equilibrium, thermal non-equilibrium, and aromatic feature emissions) of the dust emission spectrum are shown." + Starburst svstenis are iu investigateddetail i Misseltetal. (2000a)., Starburst systems are investigated in detail in \citet{mis00}. . +. We have preseuted theDIRTY radiative trauster, We have presented theDIRTY radiative transfer +elescopes detailed analysis of the shape of the line profile is not possible.,telescopes detailed analysis of the shape of the line profile is not possible. + It is possible to provide an improved fit to the data but this requires the power law of the X-ray source to rave a photon index that is significantly lower than 2., It is possible to provide an improved fit to the data but this requires the power law of the X-ray source to have a photon index that is significantly lower than 2. + We conclude that the data favour a Ixerr over a Schwarzschild Mack hole model., We conclude that the data favour a Kerr over a Schwarzschild black hole model. + Nevertheless. the work of Itevnolds Degelman (1997) ias highlighted the potential importance of [lows within the mareinally stable orbit.," Nevertheless, the work of Reynolds Begelman (1997) has highlighted the potential importance of flows within the marginally stable orbit." + Lo the irraciiating X-ray source is at some clistance above the disc. particularly in a central ocation. then fluorescent iron line [features [rom such inllowing material should be observable.," If the irradiating X-ray source is at some distance above the disc, particularly in a central location, then fluorescent iron line features from such inflowing material should be observable." + That the predicted deep absorption edge is not seen in the case of ALCG6-30-15 supports a Ixerr model and. since most of the time the line is not so broad. argues that the irradiating source lies close to the disc and. at times. changes in radius.," That the predicted deep absorption edge is not seen in the case of MCG–6-30-15 supports a Kerr model and, since most of the time the line is not so broad, argues that the irradiating source lies close to the disc and, at times, changes in radius." + The edge in the predicted spectrum of the RBOT mocel is particularly large because the material is ionized. so enhancing the contrast at the edge.," The edge in the predicted spectrum of the RB97 model is particularly large because the material is ionized, so enhancing the contrast at the edge." + lt may. be possible to have less highly ionized material within 6m if the accretion clisc were gas pressure rather than radiation pressure dominated all the way down to 67 since then the density of the disc would be many times larger., It may be possible to have less highly ionized material within $6m$ if the accretion disc were gas pressure rather than radiation pressure dominated all the way down to $6m$ since then the density of the disc would be many times larger. + Phere is also a small edge to the cold disc rellection expected in the standard model., There is also a small edge to the cold disc reflection expected in the standard model. + Lt is unlikely that such a small edge would be detectable with current X-ray telescopes. however.," It is unlikely that such a small edge would be detectable with current X-ray telescopes, however." + ]t is also useful to consider other methods of probing the innermost regions of the accretion disc., It is also useful to consider other methods of probing the innermost regions of the accretion disc. + Lf we assume rapid large variations in the continuum. are due to [Lares over the disc. then the time delay between the continuum variation and the [üorescent line response from the disc may o used to estimate the height of the Dare above the disc.," If we assume rapid large variations in the continuum are due to flares over the disc, then the time delay between the continuum variation and the fluorescent line response from the disc may be used to estimate the height of the flare above the disc." + In the case of the 14297 model this would need to be larger han that for a coronal model in order to provide sullicient illumination of the material within Gre., In the case of the RB97 model this would need to be larger than that for a coronal model in order to provide sufficient illumination of the material within $6m$. + As well as the time delay between the continuum. change ancl the response of he disc. the evolution of the iron line profile with time can ell us about the geometry of both the source and the disc.," As well as the time delay between the continuum change and the response of the disc, the evolution of the iron line profile with time can tell us about the geometry of both the source and the disc." + Unfortunately with present X-ray telescopes the photon lux in the iron line is only a few hundred counts per day. and integration times are necessarily so large that observations of short term variability are unfeasible.," Unfortunately with present X-ray telescopes the photon flux in the iron line is only a few hundred counts per day, and integration times are necessarily so large that observations of short term variability are unfeasible." + lt is exciting that we are now debating and able to distinguish. gross details of the accretion low of matter at radii less than 6m., It is exciting that we are now debating and able to distinguish gross details of the accretion flow of matter at radii less than $6m$. + Future observations with ASCA. ANAL. XMM. ASTRO-E and Constellation-N. will continue. this exploration of the very near environment of black holes.," Future observations with ASCA, AXAF, XMM, ASTRO-E and Constellation-X will continue this exploration of the very near environment of black holes." + AJY and ACE thank PPARC ancl the Roval Society. [or support. respectively.," AJY and ACF thank PPARC and the Royal Society for support, respectively." +Alore interesting is the scatter in the data.,More interesting is the scatter in the data. + For a photon counting instrument such asGALEN.. (he instrumental scatter will be either due to photon noise or to errors in (he flat fielding (calibration) of the instrument.," For a photon counting instrument such as, the instrumental scatter will be either due to photon noise or to errors in the flat fielding (calibration) of the instrument." + We have empirically derived the instrumental scatter by dividing each observation into (wo sets of visits. which may well be separated by several months.," We have empirically derived the instrumental scatter by dividing each observation into two sets of visits, which may well be separated by several months." + There is excellent agreement between this aud the intrinsic photon noise (Fig. 4)).," There is excellent agreement between this and the intrinsic photon noise (Fig. \ref{scat_plot}) )," + confirming (hat the errors are dominated by poissonian rather than instrumental effects., confirming that the errors are dominated by poissonian rather than instrumental effects. + As an independent test. we also took the overlap regions between clillerent observations and ealeulated the scatter between them.," As an independent test, we also took the overlap regions between different observations and calculated the scatter between them." + Although the seatter for the overlap regions is somewhat higher than the caleulated: values. (his is due to the many fewer points in the overlap regions and (heir location near the edge of the detector.," Although the scatter for the overlap regions is somewhat higher than the calculated values, this is due to the many fewer points in the overlap regions and their location near the edge of the detector." + We note here that all our comparisons are in skv coordinates because (here are arbitrary roll angle differences between different. visits. which do not allow a comparison between physical detector pixels.," We note here that all our comparisons are in sky coordinates because there are arbitrary roll angle differences between different visits, which do not allow a comparison between physical detector pixels." +" The FUV and NUV images of theSpitzer ""First Look” field obtained after subtraction ol the foreground emission are shown in Fig.", The FUV and NUV images of the “First Look” field obtained after subtraction of the foreground emission are shown in Fig. +" 5 at a spatial resolution of2"".", \ref{diffuse_image} at a spatial resolution of. +. The UV images of Fig., The UV images of Fig. + 5 max be compared with the IR. 100 mmap (Fie. 1))., \ref{diffuse_image} may be compared with the IR 100 map (Fig. \ref{IR_img}) ). + There are several possible contributors to the astrophysical UV. emission. a significant one being. dust-scattered starlight which contributes to both the FUV and the NUV bands.," There are several possible contributors to the astrophysical UV emission, a significant one being, dust-scattered starlight which contributes to both the FUV and the NUV bands." + This is reflected in the good correlation between the FUV and NUV bands (Fig. 6)), This is reflected in the good correlation between the FUV and NUV bands (Fig. \ref{fuv_nuv}) ) + and between the two UV bands and the IR. 100 [fluxes (Fig. 7))., and between the two UV bands and the IR 100 fluxes (Fig. \ref{UV_IR}) ). + This is in contrast with the essentially flat UW-IR curves obtained by in Region Il. The UR emission is due to thermal radiation from an optically {hin laver of dust. as the cross-section of the grains is low in the IR.," This is in contrast with the essentially flat UV-IR curves obtained by \citet{SNV09} in Region I. The IR emission is due to thermal radiation from an optically thin layer of dust, as the cross-section of the grains is low in the IR." + On the other hand. the cross-section of the grains is much higher in the UV and the optical depth transitions from being opticallv thin in these Draco observations to being optically thick in Region I. In Fig. 8..," On the other hand, the cross-section of the grains is much higher in the UV and the optical depth transitions from being optically thin in these Draco observations to being optically thick in Region I. In Fig. \ref{UV_IR_ratio}," + we have plotted the ratio between the UV bands and the IR to understaud the nature of diffuse UV emission with optical depth., we have plotted the ratio between the UV bands and the IR to understand the nature of diffuse UV emission with optical depth. + There is a clear trend. visible from the low optical depth Draco region to the hieh optical depth (in the UV) Reeion I with an empirical formula of It is interesüng to note that the ο ratio in our data follows a continuous curve verv similar to that found by Murthyetal.(2001) in Orion, There is a clear trend visible from the low optical depth Draco region to the high optical depth (in the UV) Region I with an empirical formula of It is interesting to note that the $F_{UV}/F_{IR}$ ratio in our data follows a continuous curve very similar to that found by \citet{JM01} in Orion +it has been cleared by an orbiüng giant planet. (Calvet οἱ 22002). it is tempting to inler that TW Iva shows weak molecular emission in (he mid-inliared because the planet has created a gap in thegaseous disk as well.,"it has been cleared by an orbiting giant planet (Calvet et 2002), it is tempting to infer that TW Hya shows weak molecular emission in the mid-infrared because the planet has created a gap in the disk as well." + We discussed in section 4 some of the issues (hat need to be investigated in order to determine whether the difference we observe is the result of giant planet formation or other processes., We discussed in section 4 some of the issues that need to be investigated in order to determine whether the difference we observe is the result of giant planet formation or other processes. + Given our limited understanding of the factors that govern the emission spectra of T Tauri disks. if may be useful to (take an empirical approach in exploring whether (he lack of molecular emission seen from TW Iva is a consequence of the physical evolution of the disk. its chemical evolution. or an excitation effect.," Given our limited understanding of the factors that govern the emission spectra of T Tauri disks, it may be useful to take an empirical approach in exploring whether the lack of molecular emission seen from TW Hya is a consequence of the physical evolution of the disk, its chemical evolution, or an excitation effect." + For example. to explore the possibility of an excitation effect. we can compare theSpitzer spectyum of TW Iva with those of other non-lransition objects wilh comparable accretion rates.," For example, to explore the possibility of an excitation effect, we can compare the spectrum of TW Hya with those of other non-transition objects with comparable accretion rates." + If non-(ransiGion objects with low accretion rates also lack strong molecular emission. the spectzum of TW Iva would not be unusual for its accretion rate and the lack of molecular emission max be the result of poor excitation.," If non-transition objects with low accretion rates also lack strong molecular emission, the spectrum of TW Hya would not be unusual for its accretion rate and the lack of molecular emission may be the result of poor excitation." + However. if strong molecular emission is observed in other low accretion rate svslenms. (he absence of such emission in TW Iva would suggest a gap in ils gaseous clisk or a possible chemical effect.," However, if strong molecular emission is observed in other low accretion rate systems, the absence of such emission in TW Hya would suggest a gap in its gaseous disk or a possible chemical effect." + We will take Chis approach in a future study., We will take this approach in a future study. + Sinularly. it would be interesting to explore whether other transition objects also show weak molecular emission compared (o classical T Tauri stars.," Similarly, it would be interesting to explore whether other transition objects also show weak molecular emission compared to classical T Tauri stars." +" Lf transition objects with much weaker UV fluxes than that of TW Ilva (e.g.. DM Tan. GAL Aur) also show a clelicil ol molecular emission. that would suggest a more dominant role lor SED evolution (a deficit of grains: possible clearing by a eiut planet). rather than photochemistry. in accounting for the dilference in (he spectra,"," If transition objects with much weaker UV fluxes than that of TW Hya (e.g., DM Tau, GM Aur) also show a deficit of molecular emission, that would suggest a more dominant role for SED evolution (a deficit of grains; possible clearing by a giant planet), rather than photochemistry, in accounting for the difference in the spectra." + We will report on Chis in future publications., We will report on this in future publications. + While strong molecular emüssion is not detected [rom TW Ilva. we do detect a rich spectrum of emission lines of atoms (LU. [Nel]. and [NelH]) and molecules (II5. ΟΠ. CO». HCO. and possibly ClI4).," While strong molecular emission is not detected from TW Hya, we do detect a rich spectrum of emission lines of atoms (HI, [NeII], and [NeIII]) and molecules $\Htwo,$ OH, $\COtwo,$ $\HCOp,$ and possibly $\CHthree$ )." + One of the most intriguing is (he OIL emission. which is hot and may result [rom the UV photodissociation of water.," One of the most intriguing is the OH emission, which is hot and may result from the UV photodissociation of water." + A more detailed analvsis of the OIL enission spectrum may be able to determine whether it is produced by photodissociation., A more detailed analysis of the OH emission spectrum may be able to determine whether it is produced by photodissociation. + The properties of the molecular emission Irom TW Iva. both the ΟΠΗ and other molecules. will be analvzed in greater detail in a future study.," The properties of the molecular emission from TW Hya, both the OH and other molecules, will be analyzed in greater detail in a future study." + Because we detect multiple ILE lines. we can show Chat the IHE emission from TW Iva has a recombination spectrum.," Because we detect multiple HI lines, we can show that the HI emission from TW Hya has a recombination spectrum." + In contrast to the neon emission from TW Iva. which can be well accounted for by (primarily) stellar X-ray irradiation of the disk. the physical origin of the IHE emission is difficult to identify.," In contrast to the neon emission from TW Hya, which can be well accounted for by (primarily) stellar X-ray irradiation of the disk, the physical origin of the HI emission is difficult to identify." + As discussed in section 4. magnetospheres. disk abmospheres aud/or photoevaporative Lows could plausibly contribute to the emission ancl multiple components max play a role.," As discussed in section 4, magnetospheres, disk atmospheres and/or photoevaporative flows could plausibly contribute to the emission and multiple components may play a role." + High resolution spectroscopy of the brightest ILE lines (III 7-6 and ILE 9-7) would likely provide valuable insights into the origin of the emission., High resolution spectroscopy of the brightest HI lines (HI 7-6 and HI 9-7) would likely provide valuable insights into the origin of the emission. +Modern estimates of the mean metallicity eracient in the TOO1000 pe closest to the Galactic plane (Llartkopl Yoss 1982: Yoshii et al.,Modern estimates of the mean metallicity gradient in the $\sim 700 - 1000$ pc closest to the Galactic plane (Hartkopf Yoss 1982; Yoshii et al. + LOST: Yoss et al., 1987; Yoss et al. + LOST: Sorensen Ixnude 1994: Buser Rone 1995: Trefzeer et al., 1987; rensen Knude 1994; Buser Rong 1995; Trefzger et al. + 1995: see also Robin et al., 1995; see also Robin et al. + 1996: Buser et al., 1996; Buser et al. + 1998) are in the range 0.6xdFe/H]/dz<0.3 dex *.," 1998) are in the range $-0.6 \le \hbox{\rm +d[Fe/H]/d{\it z}} \le -0.3$ dex $^{-1}$." + ME of these results have been obtained from samples. along sight lines closely perpendicular to the Galactic plane ancl centered on the position of the Sun. thereby ignoring possible radial metallicity eracients.," All of these results have been obtained from samples, along sight lines closely perpendicular to the Galactic plane and centered on the position of the Sun, thereby ignoring possible radial metallicity gradients." + Jonch-Sorensen (1995) observed FE and carly C-type (main sequence) stars in six selected directions of the Galaxy and. tried. to solve for radial and vertical eracients simultaneously., rensen (1995) observed F and early G-type (main sequence) stars in six selected directions of the Galaxy and tried to solve for radial and vertical gradients simultaneously. + His best results for z«TOO pc are 0.2E0.3 dex + and 0.01£0.08 dex + for the vertical ancl racial gradients. respectively.," His best results for $z < +700$ pc are $-0.2 \pm 0.3$ dex $^{-1}$ and $-0.01 \pm 0.03$ dex $^{-1}$ for the vertical and radial gradients, respectively." + This may indicate that the elects of such a small radial metallicity e&racient on the expected. vertical abundance gradients are small or negligible. in particular if the origin of either. eradient is clilferent.," This may indicate that the effects of such a small radial metallicity gradient on the expected vertical abundance gradients are small or negligible, in particular if the origin of either gradient is different." + Colour distributions and colour. gradients are sensitive to the metal abundances and their gradients of the integrated stellar populations in galaxies., Colour distributions and colour gradients are sensitive to the metal abundances and their gradients of the integrated stellar populations in galaxies. + Althoughradial colour gradients in moderately inclined and. face-on galaxies have been studied extensively. in general indicating blucr colours with increasing ealactocentric distance (e.g.. de Jong 1996. and references therein). only for a few relatively large. ancl well-resolvec edge-on galaxiesvertical colour gradients have been measurecl.," Although colour gradients in moderately inclined and face-on galaxies have been studied extensively, in general indicating bluer colours with increasing galactocentric distance (e.g., de Jong 1996, and references therein), only for a few relatively large and well-resolved edge-on galaxies colour gradients have been measured." + In highly inclined galaxies. the interpretation of intrinsic colours ancl colour gradients is severely. hamperec by the presence of dust in the ealactic planes.," In highly inclined galaxies, the interpretation of intrinsic colours and colour gradients is severely hampered by the presence of dust in the galactic planes." +" However. from. a comparison with published colours of moderately inclines Se galaxies. Ixuchinski ""JTerndrup (1996) have shown tha for these [ate-tvpe galaxies there is little or no reddening away [rom the cust lane."," However, from a comparison with published colours of moderately inclined Sc galaxies, Kuchinski Terndrup (1996) have shown that for these late-type galaxies there is little or no reddening away from the dust lane." + Since statistical studies have shown that the dust content of Sc galaxies is large compared. to other clise-clominatecl galaxy types (e.g. de Crijs ct al.," Since statistical studies have shown that the dust content of Sc galaxies is large compared to other disc-dominated galaxy types (e.g., de Grijs et al." + 1997). we may assume that the elfects of reddening on the intrinsic galaxy colours away from the dust lane are largest for these galaxy types.," 1997), we may assume that the effects of reddening on the intrinsic galaxy colours away from the dust lane are largest for these galaxy types." + Thus. colours ancl colour graclients measured at those distances from the galactic planes where the influence of the dust lane is negligible likely rellect theinfrinsic galactic properties.," Thus, colours and colour gradients measured at those distances from the galactic planes where the influence of the dust lane is negligible likely reflect the galactic properties." + Detailed studies of the intrinsic colours of galactic discs perpendicular to their planes (e.g.. llamabe et al.," Detailed studies of the intrinsic colours of galactic discs perpendicular to their planes (e.g., Hamabe et al." + 1979: Leevi Gerber 1979: van der Ixruit Searle 1981: Jensen Thuan 1982) are consistent with a small or no vertical colour gradient outside the cust [ane region (see also de Cirijs ct al., 1979; Hegyi Gerber 1979; van der Kruit Searle 1981; Jensen Thuan 1982) are consistent with a small or no vertical colour gradient outside the dust lane region (see also de Grijs et al. + 1997)., 1997). + Although colour gradients along the minor axis may be due to some intrinsic buleeproperty’... van der. Ixruit Searle (1981). observed. that. at. various. galactocentric distances. the vertical colours of NGC SOL are. getting systematically bluer with greater height above the plane.," Although colour gradients along the minor axis may be due to some intrinsic bulge, van der Kruit Searle (1981) observed that, at various galactocentric distances, the vertical colours of NGC 891 are getting systematically bluer with greater height above the plane." + On the other hand. Jensen Thuan (1982) did not find any evidence for a similar vertical colour eraclient in NGC 4565 in the region where the old thin disc dominates.," On the other hand, Jensen Thuan (1982) did not find any evidence for a similar vertical colour gradient in NGC 4565 in the region where the old thin disc dominates." + However. as soon as the light of the thick cise starts to dominate a small perpendicular colour gradient is present in their data. in the sense that the dise colours become redder with increasing clistance from the galactic plane.," However, as soon as the light of the thick disc starts to dominate a small perpendicular colour gradient is present in their data, in the sense that the disc colours become redder with increasing distance from the galactic plane." + A similar result has recently. been obtained. for NCC 5907 (Lequeux ct al., A similar result has recently been obtained for NGC 5907 (Lequeux et al. + 1996. 1998: Ruely et al.," 1996, 1998; Rudy et al." + 1997). which was interpreted as an extended stellar halo redder than the galactic disc or a very thick cise component.," 1997), which was interpreted as an extended stellar halo redder than the galactic disc or a very thick disc component." + The conversion of broad-banc colour gradients. to abundance ancl population gradients in external galaxies is controversial. unfortunately.," The conversion of broad-band colour gradients to abundance and population gradients in external galaxies is controversial, unfortunately." + For the cetailed analysis of the uminosity and colour profiles of edge-on galaxies one needs o adopt assumptions concerning the evolutionary stellar population svnthesis. the initial mass function. the metallicity and the star formation history. as well as about 1ο dust geometry and its eharacteristies.," For the detailed analysis of the luminosity and colour profiles of edge-on galaxies one needs to adopt assumptions concerning the evolutionary stellar population synthesis, the initial mass function, the metallicity and the star formation history, as well as about the dust geometry and its characteristics." + Due to the relative insensitivity of broad-band colours to these characteristics. in particular because of the age/metallicity degeneracy in 10 colours. of an integrated stellar population (Worthey 1994). spectral line studies seem to be a more cllective tool to isentangle metallicity and age effects. as well as population eracicnts.," Due to the relative insensitivity of broad-band colours to these characteristics, in particular because of the age/metallicity degeneracy in the colours of an integrated stellar population (Worthey 1994), spectral line studies seem to be a more effective tool to disentangle metallicity and age effects, as well as population gradients." + llowever. spectral line strength indices are relatively wud to measure. anc are also degenerate to age and metallicity. although to a lesser extent iun broad-band colours.," However, spectral line strength indices are relatively hard to measure, and are also degenerate to age and metallicity, although to a lesser extent than broad-band colours." + Sincelocal colours correlate strongly with each other (cde Jong 1996. Peleticr Balceells 1996. 1997) they can be used. as indicators of the gross properties of galaxies. in the absence of dust. where the various wavelength. ranges can oe used as diagnosties for dillerent overall galaxy properties.," Since colours correlate strongly with each other (de Jong 1996, Peletier Balcells 1996, 1997) they can be used as indicators of the gross properties of galaxies, in the absence of dust, where the various wavelength ranges can be used as diagnostics for different overall galaxy properties." + In this respect. de Jong's (1996) statistical studs is one of 1e first large surveys of spiral galaxy properties based on =uultiple passband optical ancl near-infrared observations.," In this respect, de Jong's (1996) statistical study is one of the first large surveys of spiral galaxy properties based on multiple passband optical and near-infrared observations." + Fisher. Franx Ulineworth (1996) published. one. of we very few studies dealing. with abundance gradients »pendicular to the galactic planes in. highly-inclined galaxies other than our own.," Fisher, Franx Illingworth (1996) published one of the very few studies dealing with abundance gradients perpendicular to the galactic planes in highly-inclined galaxies other than our own." + Basecl on Mg» spectral line observations of 20 S0 galaxies. they conclude that the minor axis behaviour of the 9 galaxies in their ccdge-on subsample is noticeably cillerent from that found along the major axis.," Based on $_2$ spectral line observations of 20 S0 galaxies, they conclude that the minor axis behaviour of the 9 galaxies in their edge-on subsample is noticeably different from that found along the major axis." + Whereas the major axis Meg». profiles decrease with raclius and Hatten as the bulge light contribution decreases and the disc starts to dominate. the minor axis graclicnts display a uniformly decreasing Mg» strength. with distance from the," Whereas the major axis $_2$ profiles decrease with radius and flatten as the bulge light contribution decreases and the disc starts to dominate, the minor axis gradients display a uniformly decreasing $_2$ strength with distance from the" +characterising the transition of LBVs from quiescence to outburst phases and vice-versa. in a close match to evidence gathered at more conventional (blue) optical wavelengths: (v) offered the opportunity to test quantitatively the soundness of photoronisation modelling of the rich emission line spectrum observable over the RAVE range: (vii) discovered in R 127 the presence. and quantified the physical properties of. a massive detached tontsed shell which was ejected during the 1982-2000 outburst.,"characterising the transition of LBVs from quiescence to outburst phases and vice-versa, in a close match to evidence gathered at more conventional (blue) optical wavelengths; $v$ ) offered the opportunity to test quantitatively the soundness of photoionisation modelling of the rich emission line spectrum observable over the RAVE range; $vii$ ) discovered in R 127 the presence, and quantified the physical properties of, a massive detached ionised shell which was ejected during the 1982-2000 outburst." +deposition ean be seen by all three LCs diverging at ~ 40 clavs.,deposition can be seen by all three LC's diverging at $\sim$ 40 days. + As the Ni mass is increased the secondary peak increases in both absolute magnitude and width., As the Ni mass is increased the secondary peak increases in both absolute magnitude and width. + In fact when the Ni mass is 0.14 M. (he secondary peak is brighter than the plateau. and in the V light curve the secondary peak is even brighter (han the initial peak (fig.," In fact when the Ni mass is 0.14 $_\odot$ the secondary peak is brighter than the plateau, and in the V light curve the secondary peak is even brighter than the initial peak (fig." + 106)., 10c). + At late times the V light curve have similar (ail slopes but scaled to lower Iuminositv (equation 3)., At late times the V light curve have similar tail slopes but scaled to lower luminosity (equation 3). + For comparison a moclel is shown that contains no Ni mass and abrutlv falls after the plateau., For comparison a model is shown that contains no Ni mass and abrutly falls after the plateau. + Since having no Ni mass eliminates (he secondary peak it is reasonable (o assume that Ni heating plavs a role in producing the secondary peak Cig The plateau is actually lengthened by increased amount of Ni as seen in all eraphis to the point where it is almost doubled for 0.14 M. Ni., Since having no Ni mass eliminates the secondary peak it is reasonable to assume that Ni heating plays a role in producing the secondary peak (fig The plateau is actually lengthened by increased amount of Ni as seen in all graphs to the point where it is almost doubled for 0.14 $_\odot$ Ni. + The time of the secondary. peak maximum increases with increasine Ni mass., The time of the secondary peak maximum increases with increasing Ni mass. + This can be explainecl bv the Ni-Co clecay keeping the malerial hotter for a longer time. (hus slowing the RW.," This can be explained by the Ni-Co decay keeping the material hotter for a longer time, thus slowing the RW." + Figure 10b shows the photospheric temperature for the 0.14 AL. Ni model stays hot for greater than 115 days. indicating that (he gamma rays participate in the heating of the material below the photosphere.," Figure 10b shows the photospheric temperature for the 0.14 $_\odot$ Ni model stays hot for greater than 115 days, indicating that the gamma rays participate in the heating of the material below the photosphere." + There is little change in the photopspheric velocity (fie., There is little change in the photopspheric velocity (fig. + 10d) which shows no difference until about 55 davs., 10d) which shows no difference until about 55 days. + The slieht difference in velocity. after 55 days is due to an increase in the opacity. causing the photosphere to move into faster moving material.," The slight difference in velocity after 55 days is due to an increase in the opacity, causing the photosphere to move into faster moving material." + The slopes of all the (ails (figs., The slopes of all the tails (figs. + 10a. 10c) are similar since there is no variation in mass. energy. or Ni mixing.," 10a, 10c) are similar since there is no variation in mass, energy, or Ni mixing." + This can be easily explained since all models have almost identical density. velocity. and temperature profiles fies.," This can be easily explained since all models have almost identical density, velocity, and temperature profiles figs." + 11a. Lib. lle.," 11a, 11b, 11c." + In figure Hc the heating due to Ni can be seen to slightly affect the temperature profile at day 47 in the region < 6 M..., In figure 11c the heating due to Ni can be seen to slightly affect the temperature profile at day 47 in the region $<$ 6 $_\odot$. + Figure 1d shows (he most predominate allect of Ni mass., Figure 11d shows the most predominate affect of Ni mass. + The Iuminosityv. profile directly reflects the amount of enerey supplied by racdioactivitv. similar profiles but. different. absolute Iuninosiües.," The luminosity profile directly reflects the amount of energy supplied by radioactivity, similar profiles but different absolute luminosities." + In general since the only source of energv at lates times is the decay of Co it is possible to estimate (ae Ni mass based on the absolute magnitude of the tail., In general since the only source of energy at lates times is the decay of Co it is possible to estimate the Ni mass based on the absolute magnitude of the tail. + However. a LC tail wilh a steeper slope than the decay rate of Co indicates (hal gamma ravs are escaping the ejecta and (hus an under-estimate of ihe Ni mass would be Figure 12a shows that lowering the II envelope mass enables Ni to power the full plateau.," However, a LC tail with a steeper slope than the decay rate of Co indicates that gamma rays are escaping the ejecta and thus an under-estimate of the Ni mass would be Figure 12a shows that lowering the H envelope mass enables Ni to power the full plateau." + The affects of Ni appear at day. 20 due to the recombination wave moving quickly through the low mass Il envelope and uncover the regions heated by gamma. ravs faster., The affects of Ni appear at day 20 due to the recombination wave moving quickly through the low mass H envelope and uncover the regions heated by gamma rays faster. + As the Ni mass is increased the Ie core is kept hot enough to allow the RW to move more slowly., As the Ni mass is increased the He core is kept hot enough to allow the RW to move more slowly. + At later times the tails fall faster. when compared to figure I0a. due to the faster velocity. low density and (hus less (rapping of gammia Figure 12b shows that confining the Ni to < 0.3 M. delays the affects of the Ni energy source and consequently the LC continues to [all 15 days longer than the other models in," At later times the tails fall faster, when compared to figure 10a, due to the faster velocity, low density and thus less trapping of gamma Figure 12b shows that confining the Ni to $<$ 0.3 $_\odot$ delays the affects of the Ni energy source and consequently the LC continues to fall 15 days longer than the other models in" +during outburst.,during outburst. + For the purpose of determining an upper limit to the optical Inmiinosity of the outburst. we assume (hat the counterpart candidate with 5=24.7050.08 was the counterpart.," For the purpose of determining an upper limit to the optical luminosity of the outburst, we assume that the counterpart candidate with $B=24.70\pm0.08$ was the counterpart." + With Ny=5x107 7 (see Table 4)). applving the relation of Predehl&Schmitt(1995) and a standard extinction law. we find τν=0.4.," With $_H = +5\times10^{20}$ $^{-2}$ (see Table \ref{spectab1}) ), applying the relation of \citet{predehl1995} and a standard extinction law, we find $A_B = 0.4$." + Assiuning m-M-—2447. the absolute 2 magnitude was B=—0.1740.08.," Assuming $m$ -M=24.47, the absolute $B$ magnitude was $B=-0.17\pm0.08$." + Assuming an intrinsic D—V. color of —0.09=0.14 [the mean of the Galactic LMXD catalog of Liuetal.(2001)]]. we find A=—0.08£0.16. or an upper-Hmit of Af2—0.24.," Assuming an intrinsic $B-V$ color of $-0.09\pm0.14$ [the mean of the Galactic LMXB catalog of \citet{liu2001}] ], we find $M_V=-0.08\pm0.16$, or an upper-limit of $M_V\geq-0.24$." + Our measurements allow us to provide a rough prediction of the orbital period of r1-36 using (he empirical relation between X-ray. Iuminosity. optical Iuminosityv. and orbital period for Galactie LAINBs determined by vanParadijs&MeClintoek.(1994).," Our measurements allow us to provide a rough prediction of the orbital period of r1-36 using the empirical relation between X-ray luminosity, optical luminosity, and orbital period for Galactic LMXBs determined by \citet{vanparadijs1994}." +.. This relation has been tested for more recent iransient events and lor observations separated by up (o 3 weeks bv Williamsοἱal.(2005a.d).. showing that it provides reliable orbital period predictions even for events with complex mulliwaveleneth lishteurves.," This relation has been tested for more recent transient events and for observations separated by up to 3 weeks by \citet{williams2005bh1,williams2005bh4}, showing that it provides reliable orbital period predictions even for events with complex multiwavelength lightcurves." + In addition. since r1-36 shows a decay curve reminiscent of the second-half of the 1998-1999 outburst of NTE J1550-564. we tested the relation lor this ease.," In addition, since r1-36 shows a decay curve reminiscent of the second-half of the 1998-1999 outburst of XTE J1550-564, we tested the relation for this case." + The second half of that outburst had an X-ray luminosity of ~2x105 erg ! (for a distance of 5.342.3 kpe: Oroszetal. 2002)). and the optical counterpart showed Wo~16.3.," The second half of that outburst had an X-ray luminosity of $\sim2\times 10^{38}$ erg $^{-1}$ (for a distance of $\pm$ 2.3 kpc; \citealp{orosz2002j1550}) ), and the optical counterpart showed $V\sim16.3$." +" Applving the extinction of ly=4.75 (Oroszetal.2002) gives My=—2.d""d", Applying the extinction of $A_V=4.75$ \citep{orosz2002j1550} gives $_V = -2.1^{+1.2}_{-0.8}$. + These numbers would vield a period prediction of 2 1 dax., These numbers would yield a period prediction of $\gap$ 1 day. + This limit is correct. as (the (rue period is 1.55 days.," This limit is correct, as the true period is 1.55 days." + In the case of r1-36. if we insert the optical Iuminositv of Ady=—0.08zc0.16 and the unabsorbed 0.3.7 keV X-ray luminosity of 5x LO ere ! into the van relation. we obtain a prediction for the orbital period of 1.77) davs.," In the case of r1-36, if we insert the optical luminosity of $M_V=-0.08\pm0.16$ and the unabsorbed 0.3–7 keV X-ray luminosity of $\times$ $^{37}$ erg $^{-1}$ into the \citet{vanparadijs1994} + relation, we obtain a prediction for the orbital period of $^{+3.5}_{-1.0}$ days." + Therefore. assuming r1-36 is an LMXD similar to those in our own Galaxy and that the true counterpart was no brighter (han the only star detected within the 3e error ellipse during the outburst. the predicted upper-Iimit for the period of the svstem is Z 5.2 days.," Therefore, assuming r1-36 is an LMXB similar to those in our own Galaxy and that the true counterpart was no brighter than the only star detected within the $\sigma$ error ellipse during the outburst, the predicted upper-limit for the period of the system is $\lap$ 5.2 days." + Ideally we would like to fully elassify the transient N-ravy source r1-36., Ideally we would like to fully classify the transient X-ray source r1-36. + While its observed properties show that it is an LMXD in MBI. the current. observations do not allow a final conclusion to be drawn as to whether the primary member of (he binary is a neutron star or black hole.," While its observed properties show that it is an LMXB in M31, the current observations do not allow a final conclusion to be drawn as to whether the primary member of the binary is a neutron star or black hole." + Nevertheless. the similarities between r1-36 and X-ray binaries known (o contain black holes make it a good black hole candidate.," Nevertheless, the similarities between r1-36 and X-ray binaries known to contain black holes make it a good black hole candidate." +(Goldreich&Lynden-Dell1965).,\citep{GoldreichLyndenBell65}. +. The source terms in (his approximation can be dropped to study other problems., The source terms in this approximation can be dropped to study other problems. + We choose a local reference frame located at a fiducial radius. corotating at the orbital angular velocity Q.," We choose a local reference frame located at a fiducial radius, corotating at the orbital angular velocity $\Omega$." + The clvnamical equations are written using the Carlesian coordinate. wilh ag.z denoting unit vectors pointing to the radial. azimuthal and vertical direction.," The dynamical equations are written using the Cartesian coordinate, with $\hat{\mb{x}},\hat{\mb{y}},\hat{\mb{z}}$ denoting unit vectors pointing to the radial, azimuthal and vertical direction." +" In (his non-inertial lraune. the coupled equations of particles and gas read In the above equations. py. 2, denote the mass density anc pressure of the gas. dlnO/dlnris the backerounc shear parameter. wilh q=3/2 lor Ixeplerian flow. ancl a. v denote velocities of the gas and particles in this reference Tame."," In this non-inertial frame, the coupled equations of particles and gas read In the above equations, $\rho_g$, $P_g$ denote the mass density and pressure of the gas, $q\equiv d\ln\Omega/d\ln r$ is the background shear parameter, with $q=3/2$ for Keplerian flow, and $\mb{u}$, $\mb{v}$ denote velocities of the gas and particles in this reference frame." + The subseript πο in equation (1)) represents the 7th particle.," The subscript $i$ "" in equation \ref{eq:dustmotion}) ) represents the $i$ th particle." + The particle stopping (ime due to gas drag. lop: depends on particle size aud gas [low properties (Weidenschiliing1977).," The particle stopping time due to gas drag, $t_{\rm stop}$, depends on particle size and gas flow properties \citep{Weidenschilling77}." +. Our code is capable of dealing with an arbitrary number of different particle species (each particle species has a different stopping time). but unless otherwise stated. we assume single particle species with constant stopping time throughout this paper for simplicitv.," Our code is capable of dealing with an arbitrary number of different particle species (each particle species has a different stopping time), but unless otherwise stated, we assume single particle species with constant stopping time throughout this paper for simplicity." +" ln equation (3)). v stands for averaged particle velocity in (he ""Iud element” (weiehted by mass). and e denotes the local nass density ratio between particle and gas €=pyj/p,."," In equation \ref{eq:gasmotion}) ), $\overline{\mb{v}}$ stands for averaged particle velocity in the “fluid element"" (weighted by mass), and $\epsilon$ denotes the local mass density ratio between particle and gas $\epsilon=\rho_p/\rho_g$." + This term represents momentum feedback from the particles to (he gas. written in the form of treating particles as a fluid.," This term represents momentum feedback from the particles to the gas, written in the form of treating particles as a fluid." + The particle treatinent of feedback term is described in relssec:scheme.. and conservation of total momentum is guaranteed.," The particle treatment of feedback term is described in \\ref{ssec:scheme}, and conservation of total momentum is guaranteed." + In tliis paper. we consider non-stratified disks bv neglecting vertical gravity terms in the equations above (i.e.. the 07: terms).," In this paper, we consider non-stratified disks by neglecting vertical gravity terms in the equations above (i.e., the $\Omega^2z$ terms)." + We also neglect terms associated with the magnetic field in this paper., We also neglect terms associated with the magnetic field in this paper. + Thev are handled bv the underlying MIID integrators in Athena (Stoneetal.2008).., They are handled by the underlying MHD integrators in Athena \citep{AthenaTech}. + An isothermal equation. of. state for. the gas is. used throughout this. paper. withB P—>pc2.," An isothermal equation of state for the gas is used throughout this paper, with $P=\rho_gc_s^2$." + Our goal is to perform the local shearing box simulations (Llawlevetal.1995).. where the radial boundary. condition is periodic with additional shear to account [or differential rotation.," Our goal is to perform the local shearing box simulations \citep{HGB95}, where the radial boundary condition is periodic with additional shear to account for differential rotation." + Therefore. it is not appropriate to include radial pressure gradient directly. which is inconsistent with the periodie boundary conditions.," Therefore, it is not appropriate to include radial pressure gradient directly, which is inconsistent with the periodic boundary conditions." + Alternativelv. one can replace the pressure gradient by à constant radial force acting on the gas E—2eyGa. poimüngoulware.," Alternatively, one can replace the pressure gradient by a constant radial force acting on the gas ${\mb F}=2\eta v_K\Omega\hat{\mb x}$, pointing." + The quantity egy measures the amount by which the gas (azimuthal) velocity is reduced from (he Keplerian value due to (he radial pressure gradient., The quantity $\eta v_K$ measures the amount by which the gas (azimuthal) velocity is reduced from the Keplerian value due to the radial pressure gradient. + In our code. instead. we find it more," In our code, instead, we find it more" +which gave the shock properties such as the shock stanclolf distance A. the Mach number AS. and (he radius of curvature at the nose of the obstacle Ao.,"which gave the shock properties such as the shock standoff distance $\Delta$, the Mach number $M$, and the radius of curvature at the nose of the obstacle $R_{O}$." + Figure 3((a) shows data and the initial fit. Figure 3((b) shows the shifted data and the shock fit using Equation (5)).," Figure \ref{f3}( (a) shows data and the initial fit, Figure \ref{f3}( (b) shows the shifted data and the shock fit using Equation \ref{Eq6}) )." +" The fast maegnetosonic Mach number was calculated using Mj,=(6,4,—Ose)μι. where v4, is the CME velocity. c4, is the solar wind velocity and 0,,, is the fast magnetosonic speed."," The fast magnetosonic Mach number was calculated using $M_{ms} = ({v_{cme}-v_{sw}})/{v_{ms}}$, where $v_{cme}$ is the CME velocity, $v_{sw}$ is the solar wind velocity and $v_{ms}$ is the fast magnetosonic speed." +" Since Cae and 04,; were not known at the position of the CALE. a model corona was used to evaluate them."," Since $v_{sw}$ and $v_{ms}$ were not known at the position of the CME, a model corona was used to evaluate them." + This was based on the Parker solar wind solution with a simple dipolar magnetic field of the form Bir)=By(R./r)*. where was GG at the solar surface (Mannetal.2003)..," This was based on the Parker solar wind solution with a simple dipolar magnetic field of the form $B(r)=B_{0}(R_{\odot}/r)^{3}$, where was G at the solar surface \citep{Mann:2003p9016}." +" For each of the paired CME and shock observations the standoll distances A (=D,—Do) were obtained by three dillerent means: (1) using the 3D coordinates of the furthest point (max(/). where fh=y?+i? 27) on the shock and the CME as fy, and hi, respectively. Gi) the previous method can be applied but to the data in the common coordinate svstem which gave De and Dy. ancl (11) the Iront fitting procedure also produced stancoll distauces."," For each of the paired CME and shock observations the standoff distances $\Delta$ $D_{S}-D_{O}$ ) were obtained by three different means: (i) using the 3D coordinates of the furthest point $\textrm{max}(h)$, where $h=\sqrt{x^2+y^2+z^2}$ ) on the shock and the CME as $h_{shk}$ and $h_{cme}$ respectively, (ii) the previous method can be applied but to the data in the common coordinate system which gave $D_{O}$ and $D_{S}$, and (iii) the front fitting procedure also produced standoff distances." + llowever the results of method (1) cannot be used with the relations from Section 1. as they are not ina CME/obstacle centred coordinate svstem. but (the results [rom method (ii) and ii) can be compared to Equations (2)). (2)) and (4)).," However the results of method (i) cannot be used with the relations from Section \ref{s_intro} as they are not in a CME/obstacle centred coordinate system, but the results from method (ii) and (iii) can be compared to Equations \ref{Eq2}) ), \ref{Eq4}) ) and \ref{Eq5}) )." + A summary of the shock properties derived [rom the observations is shown in Figure 4((aj-() as a function of me., A summary of the shock properties derived from the observations is shown in Figure \ref{f4}( (a)-(f) as a function of time. +" With the exception of the CME (/,,,) and shock heights (Pha). All Che properties have been derived [rom the data collapsed on to a common coordinate svslem wilh respect to the CME."," With the exception of the CME $h_{cme}$ ) and shock heights $h_{shk}$ ), all the properties have been derived from the data collapsed on to a common coordinate system with respect to the CME." + The gap between the first three data points and others is a result of both the CME and shock leaving the COR?2 [ield-of-view and entering the III field-ol-view., The gap between the first three data points and others is a result of both the CME and shock leaving the COR2 field-of-view and entering the HI1 field-of-view. + The contrast between shock and background in the first three and last three observation is extremely low. making identification of the shock difficult.," The contrast between shock and background in the first three and last three observation is extremely low, making identification of the shock difficult." + As a result. these points are not reliable. aud shoulel be neglected.," As a result, these points are not reliable, and should be neglected." + Figure 4((a) shows (he derived heights of the CMEancl shock as they were tracked from SRR. to RIL. AAU)., Figure \ref{f4}( (a) shows the derived heights of the CMEand shock as they were tracked from $_{\odot}$ to $_{\odot}$ AU). + Using a linear fil 10 hag—hus (=A) versus hi. (not shown). the extrapolated stancolf distance at Earth was Iound to be ~4O RI...," Using a linear fit to $h_{shk}-h_{cme}$ $\Delta$ ) versus $h_{cme}$ (not shown), the extrapolated standoff distance at Earth was found to be $\sim$ $_{\odot}$." + Figure 4((b) shows the distance to nose of the CME (De) and shock (Ds) front by (filled svanbols). also shown are the values derived [rom fits to the shock aud CME front (hollow symbols).," Figure \ref{f4}( (b) shows the distance to nose of the CME $D_{O}$ ) and shock $D_{S}$ ) front by (filled symbols), also shown are the values derived from fits to the shock and CME front (hollow symbols)." + The increasing offset between the (wo is due to their differing centres of the coordinate svstems. as one is elliptic and the other is parabolic.," The increasing offset between the two is due to their differing centres of the coordinate systems, as one is elliptic and the other is parabolic." + Figure 4((¢) shows the stanclolf distance A derived using Do and Dy (fillel symbols) and from the fits to the fronts (hollow svimbols)., Figure \ref{f4}( (c) shows the standoff distance $\Delta$ derived using $D_{O}$ and $D_{S}$ (filled symbols) and from the fits to the fronts (hollow symbols). + Both are in general agreement and show an increase with time., Both are in general agreement and show an increase with time. + The standolf distance normalised using De is shown in Figure ία}., The standoff distance normalised using $D_{O}$ is shown in Figure \ref{f4}( (d). + The normalised stancloll distance is roughly constant with à mean value of 0.37 £0.09., The normalised standoff distance is roughly constant with a mean value of $0.37\pm0.09$ . + The stancloll distance, The standoff distance + ~300 2=1 ;~2 ;~0. o CA) , $\sim 300$ $z > 1$ $z \sim 2$ $z \sim 0$ $\sigma$ $A$ +Weak lensing provides a unique method to directly measure the mass lluctuations on large scales in the universe (see Alellier 1999: Ixaiser 1999: Dartelmann Schneider 1999 [or recent reviews).,Weak lensing provides a unique method to directly measure the mass fluctuations on large scales in the universe (see Mellier 1999; Kaiser 1999; Bartelmann Schneider 1999 for recent reviews). +" ""This method relies on the measurement of small. coherent. distortions produced. by lensing upon the shapes of background galaxies."," This method relies on the measurement of small, coherent distortions produced by lensing upon the shapes of background galaxies." + This elect is now routinely used to map the mass of clusters of galaxies (see reviews by Fort Alellier 1994. Schneider 1996).," This effect is now routinely used to map the mass of clusters of galaxies (see reviews by Fort Mellier 1994, Schneider 1996)." + ltecently. the technique was extended to the field by several groups who reported the statistical detection of weak lensing by large-scale structure (Wittman et al.," Recently, the technique was extended to the field by several groups who reported the statistical detection of weak lensing by large-scale structure (Wittman et al." + 2000: van Waerbeke ot al., 2000; van Waerbeke et al. + 2000: )acon. Itefregier Elis 2000 (BRE): Ixaiser. Wilson Luppino 2000).," 2000; Bacon, Refregier Ellis 2000 (BRE); Kaiser, Wilson Luppino 2000)." +" More precise measurements of this ""cosmic shear” from upcoming observations will provide invaluable cosmological information (eg.", More precise measurements of this “cosmic shear” from upcoming observations will provide invaluable cosmological information (eg. + Ixaiser 1992: Jain Seljak 1997: Ixamionkowski et al., Kaiser 1992; Jain Seljak 1997; Kamionkowski et al. + LOOT: Ixaiser 1998: Hu Tegmark 1998: van Waerbeke et al., 1997; Kaiser 1998; Hu Tegmark 1998; van Waerbeke et al. + 1998)., 1998). + Because the distortions induced by lensing are only of the order of114.. these measurements are very challenging.," Because the distortions induced by lensing are only of the order of, these measurements are very challenging." + In particular. they require tight control of systematic effects and a precise method for the measurement of the shear.," In particular, they require tight control of systematic effects and a precise method for the measurement of the shear." + One of the potential weaknesses of the cosmic shear programme, One of the potential weaknesses of the cosmic shear programme +svslenmatic errors actually dominate in all emission models. which is tvpical for modeling of astrophysical observations of relatively complex phenomena.,"systematic errors actually dominate in all emission models, which is typical for modeling of astrophysical observations of relatively complex phenomena." + The X-ray and TeV 5-rav data contribute the most to the overall 47., The X-ray and TeV $\gamma$ -ray data contribute the most to the overall $\chi^2$. + Especially for the TeV 5-ray. data. the 4? value is 149 for 27 data points. corresponding to an average residuals about. 2.30.," Especially for the TeV $\gamma$ -ray data, the $\chi^2$ value is $149$ for $27$ data points, corresponding to an average residuals about $2.3\sigma$." + That is to sav this simple leptonic model actually can not fit the £55 data well., That is to say this simple leptonic model actually can not fit the $HESS$ data well. + This is a well-known result in previous studies (e.g..Aharonianetal.2006;Tanaka2008:Morlino20092:Fangetal. 2009)..," This is a well-known result in previous studies \citep[e.g.,][] +{2006A&A...449..223A,2008ApJ...685..988T,2009MNRAS.392..240M, +2009MNRAS.392..925F}." +" In Linetal.(2008) the authors proposed a stochastic acceleration model io generate (he electron spectrum. with sub-exponential eutoll (0,—0.5) to better fit the ILESS data.", In \cite{2008ApJ...683L.163L} the authors proposed a stochastic acceleration model to generate the electron spectrum with sub-exponential cutoff $\delta_e=0.5$ ) to better fit the $HESS$ data. + ILowever. in such a case the fit to N-rav data becomes worse.," However, in such a case the fit to X-ray data becomes worse." +" The X-ray data actually favors super-exponential cutoff instead (with 9,=1.2 in this purely leptonic Lit).", The X-ray data actually favors super-exponential cutoff instead (with $\delta_e=1.2$ in this purely leptonic fit). + The fit max be improved in some detailed leptonic models. as shown in Fanetal.(2010b).. though svstematic errors still dominate.," The fit may be improved in some detailed leptonic models, as shown in \cite{2010A&A...517L...4F}, though systematic errors still dominate." + In (his subsection we discuss the model with a predominantly hadronie origin of the 5-, In this subsection we discuss the model with a predominantly hadronic origin of the $\gamma$ -rays. +" The spectrum of the accelerated protons is assunied to be F,U7)xE""rexp|l-(E/E?y] wilh 9,=1. which gives acceptable fit to the TeV data."," The spectrum of the accelerated protons is assumed to be $F_p(E)\propto E^{-\alpha_p}\exp[-(E/E_c^p)^{\delta_p}]$ with $\delta_p=1$, which gives acceptable fit to the TeV data." + The normalization is fixed using the total. kinetic energv of protons with the enerey E>1 GeV. For the hadronic οτανproduction we adopt the parameterization of Kamaeetal. (2006)..., The normalization is fixed using the total kinetic energy of protons with the energy $E>1$ GeV. For the hadronic $\gamma$ -rayproduction we adopt the parameterization of \cite{2006ApJ...647..692K}. . + With the additional, With the additional +lio various masses.,to various masses. + Figure 3 gives the number of black holes as a [function of time for various choices of the initial black hole mass. μμ. wilh other parameters held constant.," Figure 3 gives the number of black holes as a function of time for various choices of the initial black hole mass, $M_{BH,0}$, with other parameters held constant." + We can now use the same framework to estimate the huminositv of the ensemble of accreting stellar-mass black holes at the epoch. (.. We will adopt the parameterization of Park&RicotGi(2011) with the assumption that there is sufficient angular momentum {ο form a disk near the black hole. so that clisk-like efficiencies lor turning mass accretion rates into radiated energy are applicable.," We can now use the same framework to estimate the luminosity of the ensemble of accreting stellar-mass black holes at the epoch, t. We will adopt the parameterization of \citet{PR11} with the assumption that there is sufficient angular momentum to form a disk near the black hole, so that disk-like efficiencies for turning mass accretion rates into radiated energy are applicable." + For critiques of Chis assumption. see Iuulfert&Arnett(1994):Beskin&Ixarpov (2005).," For critiques of this assumption, see \citet{RA94,BK05}." +. We write for the luminosity of a single accreting black hole: with 5~Q.1., We write for the luminosity of a single accreting black hole: with $\eta \sim 0.1$. + If a disk does not form. or forms only sporaclically. the radiation efficiency would be correspondinglv less than the fiducial value we assume here.," If a disk does not form, or forms only sporadically, the radiation efficiency would be correspondingly less than the fiducial value we assume here." +" The luminosity per unit galaxv mass radiated by all black holes with mass between May aud Mp+dMg; al epoch. t. is given by: Using Equation 29.. (his can be written as: The total luminosity from all the accreting black holes born since /, can then be obtained bv integratingexe over all the current masses al epoch. t. to obtain: Once again. we can approximate the complex variation of the rate of production of black holes with a constant to obtain: Neglecting {η and /4,,4. taking /< 20\,$ kpc." + Therefore. we are in the »eculiar situation where the mass profile of the Galaxy. is ess well determined than for some nearby. spiral galaxies.," Therefore, we are in the peculiar situation where the mass profile of the Galaxy is less well determined than for some nearby spiral galaxies." + There is no possibility of substantially increasing the number of known satellite galaxies and globular clusters ancl the »st prospect for improving the measurement of the mass is through isolating large numbers of another distant. haloracer., There is no possibility of substantially increasing the number of known satellite galaxies and globular clusters and the best prospect for improving the measurement of the mass is through isolating large numbers of another distant halo--tracer. + Fielel blue horizontal branch. (BLIB) stars. should »ovide just such a tracer and WISZ99 ealeulate that to reduce he uncertainty on the total mass to 20% requires a sample of 200 distant. BILD stars., Field blue horizontal branch (BHB) stars should provide just such a tracer and WE99 calculate that to reduce the uncertainty on the total mass to $20\%$ requires a sample of 200 distant BHB stars. + This is the first in a series of three papers presenting a new calculation of the mass of the Galaxy using racial velocities of DIID. stars., This is the first in a series of three papers presenting a new calculation of the mass of the Galaxy using radial velocities of BHB stars. + Field BIB stars are. luminous standard candles that are abundant in the Galactic halo (c.g. Yanny et al., Field BHB stars are luminous standard candles that are abundant in the Galactic halo (e.g. Yanny et al. + 2000). and for nearly twenty vears. since the studs of Pier (1983). have presented an uncerexploited resource with which to measure the density profile and phase space structure of the Galaxy halo out to large distances. 100 kpe.," 2000), and for nearly twenty years, since the study of Pier (1983), have presented an under–exploited resource with which to measure the density profile and phase space structure of the Galaxy halo out to large distances, $\sim 100\,$ kpc." + A number of dynamical analyses of rather small samples of DIID. stars have been published (e.g. SommerLarsen. Christensen Carter. 1950. Norris Llawkins. 1991. Arnold Gilmore. 1992).," A number of dynamical analyses of rather small samples of BHB stars have been published (e.g. Sommer--Larsen, Christensen Carter, 1989, Norris Hawkins, 1991, Arnold Gilmore, 1992)." + Unfortunately. samples of field Aotype stars in the halo include not only DIID stars but also stars of main sequence surface gravity. field blue stragelers. tha are some 2 magnitudes less luminous.," Unfortunately, samples of field A–type stars in the halo include not only BHB stars but also stars of main sequence surface gravity, field blue stragglers, that are some 2 magnitudes less luminous." + Progress towards the goal of acquiring a large sample of distant. DII. stars has been slow because of the dilliculty of separating ou the BIB stars without the investment of large amounts of telescope time., Progress towards the goal of acquiring a large sample of distant BHB stars has been slow because of the difficulty of separating out the BHB stars without the investment of large amounts of telescope time. + In this first paper we describe our procedures [or classifving samples of halo Atype stars. and. presen a new efficient. method. that requires. only. spectroscopic observations of intermediate signaltonoise ratio.," In this first paper we describe our procedures for classifying samples of halo A–type stars, and present a new efficient method that requires only spectroscopic observations of intermediate signal–to–noise ratio." + In Paper Η we will present. photometry ancl spectroscopy of. fain 16«D20 candidate BIB stars in two northern high Galactic latitude fields and four southern fields., In Paper II we will present photometry and spectroscopy of faint $16 is a complex shear term. represeuting the anisotropic part of the distortion. with 5=τσ]τι aud the real and imaginary parts beiug denoted with the subscripts. ""1 and 727 respectively. as per convention."," The term $\gamma$ is a complex shear term, representing the anisotropic part of the distortion, with $\gamma=|\gamma|e^{2i\phi}$, and the real and imaginary parts being denoted with the subscripts, “1” and “2” respectively, as per convention." + Using our locally defined coordinate svsteni we have: Likewise. asstuning the absence of a caustic crossing (Low[x] <0). this can be inverted uniquely to give a projection fiction. Iu this analvsis. we focus ou expanding this projection operator to the next higher order to derive the octopole moment rather than restricting ourselves to the quadrupole moment alone.," Using our locally defined coordinate system, we have: Likewise, assuming the absence of a caustic crossing $1-\kappa-|\gamma| < 0$ ), this can be inverted uniquely to give a projection function, In this analysis, we focus on expanding this projection operator to the next higher order to derive the octopole moment rather than restricting ourselves to the quadrupole moment alone." + Iu general. researchers have treated weak lensing fields iu the manner described by NSB or its variants (Bacon. Refreeier Ellis. 2000: van Waerbeke et al.," In general, researchers have treated weak lensing fields in the manner described by KSB or its variants (Bacon, Refregier Ellis, 2000; van Waerbeke et al." + 2001)., 2001). + These tecliniques describe the mapping of source-plane quadrupole light distributions to leus-plaue distributions. aud thus use the observed ellipticity aud an assumption of random oricutation to invert the shear field.," These techniques describe the mapping of source-plane quadrupole light distributions to lens-plane distributions, and thus use the observed ellipticity and an assumption of random orientation to invert the shear field." + Du this work. we aim to generalize these transformations to the next higher order.," In this work, we aim to generalize these transformations to the next higher order." + Our notation for the n-th order moments of a galaxy is:, Our notation for the $n$ -th order moments of a galaxy is: +mediuu has been speculated to be responsible for a number of other peculiar star-forming svstenus (6.9. Beaulieu et al.,"medium has been speculated to be responsible for a number of other peculiar star-forming systems (e.g., Beaulieu et al." + 2010)., 2010). + Alternatively. the eas reservoir could have deen built slowly through multiple accretions of simaller eas clouds or streams. which could be either neutral or ionized.," Alternatively, the gas reservoir could have been built slowly through multiple accretions of smaller gas clouds or streams, which could be either neutral or ionized." + Under this scenario. star formation would need to be suppressed somehow caring the buildup of the reservoir.," Under this scenario, star formation would need to be suppressed somehow during the buildup of the reservoir." + The galaxy. formation models of Birnboim et al. (, The galaxy formation models of Birnboim et al. ( +2007) exhibit quiescent. reservoir building periods simular to what would be needed here. but in general they apply to somewhat higher mass galaxies than GASS35981. aud also may not be valid at 2~0.,"2007) exhibit quiescent, reservoir building periods similar to what would be needed here, but in general they apply to somewhat higher mass galaxies than GASS35981, and also may not be valid at $z\sim0$." + Multiple ninor mergers with sasadch dwarfs could also supplv the eas. but it becomes even harder in this case to imagine how the IIT could build up over time rather than form stars with each new accretion event.," Multiple minor mergers with gas-rich dwarfs could also supply the gas, but it becomes even harder in this case to imagine how the HI could build up over time rather than form stars with each new accretion event." +" We have reported ou the remarkable galaxy CASS35981. a disk galaxw with stellar ass =2«104) AL, which coutains an additional 2.1].«1019 MAL. of ΤΠ eas."," We have reported on the remarkable galaxy GASS35981, a disk galaxy with stellar mass $_*=2\times10^{10}$ $_\odot$ which contains an additional $2.1\times10^{10}$ $_\odot$ of HI gas." + Millimeter observations indicate a molecular eas mass only a teuth this high., Millimeter observations indicate a molecular gas mass only a tenth this high. + Through follow-up loue-slit spectroscopy. plus SED fitting using our UV through optical photometry we have shown that: The main conclusion from our observations ids that CASS35981 appears to he in the carly stages of formation ofits outer stellar disk.," Through follow-up long-slit spectroscopy, plus SED fitting using our UV through optical photometry we have shown that: The main conclusion from our observations is that GASS35981 appears to be in the early stages of formation of its outer stellar disk." + We are not able to provide conclusive answers to questions portainiug to the origin aud fate of the eas in this galaxw with this data set alone., We are not able to provide conclusive answers to questions pertaining to the origin and fate of the gas in this galaxy with this data set alone. + Scenarios iu which the eas was acquired iu a recent mereing event are cisfavoured because of the extremely regular kiuneuiatics of the disk., Scenarios in which the gas was acquired in a recent merging event are disfavoured because of the extremely regular kinematics of the disk. + The ΤΗ mass of GÀSS35981 is too huge to be casily explained by gas transfer frou a passiug ealaxy., The HI mass of GASS35981 is too large to be easily explained by gas transfer from a passing galaxy. + We therefore speculate that CCASS35981 acquired its eas directly from the interealactie medimn., We therefore speculate that GASS35981 acquired its gas directly from the intergalactic medium. + Although our observations show that the stars in the outer disk formed within the last Cir. this does mot mean that the gas was also acquired less than 1 Gyr ago.," Although our observations show that the stars in the outer disk formed within the last Gyr, this does not mean that the gas was also acquired less than 1 Gyr ago." + It is also nuclear whether CASS35981 will continue forming stars in its current low-cfhiciency state. or whether the eas will flow inwards towards the bulge. aud GASS35981 will eventually develop ito a more normal massive spiral ealaxy with a star formation surface deusity that decreases as a function of radius.," It is also unclear whether GASS35981 will continue forming stars in its current low-efficiency state, or whether the gas will flow inwards towards the bulge, and GASS35981 will eventually develop into a more normal massive spiral galaxy with a star formation surface density that decreases as a function of radius." + Questions concerning the eventual fate of the σας can be addressed using the larger samples that will be provided by the full CASS aud COLD GÀSS surveys in the future., Questions concerning the eventual fate of the gas can be addressed using the larger samples that will be provided by the full GASS and COLD GASS surveys in the future. + By studyiug trends iu SFR surface deusity. mean stellar age. metalliitv. aud stellar mass profiles as a function of atomic and molecular eas content for complete samples of galaxies. we hope to map out evolutionary sequences in disk galaxy. formation.," By studying trends in SFR surface density, mean stellar age, metallicity, and stellar mass profiles as a function of atomic and molecular gas content for complete samples of galaxies, we hope to map out evolutionary sequences in disk galaxy formation." + Answers to questions concerning the origin of the eas will likely require a different approach., Answers to questions concerning the origin of the gas will likely require a different approach. + Our comparison of the TT linescidth of CASS35981 with its CO line width and Wa rotation curve vield tantalizing hints that the atomic gas may not be iu equilixium with the rest of the ealaxy., Our comparison of the HI linewidth of GASS35981 with its CO line width and $\alpha$ rotation curve yield tantalizing hints that the atomic gas may not be in equilibrium with the rest of the galaxy. + In addition. the ΤΗ spectiuui in Figure 1 is clearly asvuunetric about the line center.," In addition, the HI spectrum in Figure \ref{gas_prof} is clearly asymmetric about the line center." + Hieh resolution III mapping of CASS35981 will be needed to uuderstiuid the dvnaiical state of the eas in more detail., High resolution HI mapping of GASS35981 will be needed to understand the dynamical state of the gas in more detail. + Even so. such observations are uulikelv to prove that the III originated ποια a inore diffuse (and unuseen) reservolr of IGAL eas.," Even so, such observations are unlikely to prove that the HI originated from a more diffuse (and unseen) reservoir of IGM gas." + This can only be doue if we are able to find tracers of this eas. for example absorption lines in the spectra of background quasars that arise when the quasar light passes through the circuimealactic medi of the ealaxy (Cou Ostriker 1999).," This can only be done if we are able to find tracers of this gas, for example absorption lines in the spectra of background quasars that arise when the quasar light passes through the circumgalactic medium of the galaxy (Cen Ostriker 1999)." + These tracers must then be linked with galaxies like GASS35981., These tracers must then be linked with galaxies like GASS35981. + The authors thauk J. Brinchinanu and €. Tremonti for making available their code for analysis of spectra., The authors thank J. Brinchmann and C. Tremonti for making available their code for analysis of spectra. + Based ou observations carried out with the IRAM 30m. telescope., Based on observations carried out with the IRAM 30m telescope. + IRAALD is supported by INSU/CNRS (France). MPG. (Cermany) and IGN (Spain).," IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain)." + The Arecibo Observatory is part of the National Astronomy and Touosphere Center. which is operated by Cornell University under a cooperative aerecient with the National Scicuce Foundation.," The Arecibo Observatory is part of the National Astronomy and Ionosphere Center, which is operated by Cornell University under a cooperative agreement with the National Science Foundation." + Observations reported here were obtained iu part at the AIMIT Observatory. a facility operated jointly by the Siuithsonian Iustitution aud the University of Arizona.," Observations reported here were obtained in part at the MMT Observatory, a facility operated jointly by the Smithsonian Institution and the University of Arizona." + MAIT telescope time was eranted by NOAQO. through the Telescope System Tustrmuentation Program (TSIP).," MMT telescope time was granted by NOAO, through the Telescope System Instrumentation Program (TSIP)." + TSIP is funded by NSF., TSIP is funded by NSF. + Funding for the SDSS has been provided bx the Alfred P. Sloan Foundation. the Participating Tustitutions. the National Science Foundation. the U.S. Department of Enerev. the National Acronautics aud Space Acuninistration. the Japanese \loubukagalasho. the Max Planck Society. and the IHigher. Education Funding Council for Euglaud.," Funding for the SDSS has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England." +"Table 7. shows the median. fifth. and 95!"" percentiles for the expected TTV signal aud SNR for each of the candidate planets.","Table \ref{ttvtable} shows the median, fifth, and $^{\text{th}}$ percentiles for the expected TTV signal and SNR for each of the candidate planets." + [istoerams sinular to figures 9 and 10 for cach svsteni are found in the Appendix., Histograms similar to figures \ref{ttv896} and \ref{ecc896} for each system are found in the Appendix. + With the possible exception of KOI 191. each of the planets in these five svstenis will likely have observable transit inning variations by the eud of au extended missiou.," With the possible exception of KOI 191, each of the planets in these five systems will likely have observable transit timing variations by the end of an extended mission." + For IKOI 191. even if the TTV signal is πα it may vet be detectable simply because here will be a laree number of transits Cuore han 1000) over the duration of au extended wission Which may compensate for the low sigual-o-hnoise ratio of the TTY signal to the transit nue uncertainties.," For KOI 191, even if the TTV signal is small it may yet be detectable simply because there will be a large number of transits (more than 1000) over the duration of an extended mission which may compensate for the low signal-to-noise ratio of the TTV signal to the transit time uncertainties." + The primary reason for the small signal iu KOI 191 is the large ratio of orbital periods. exceeding 6:1.," The primary reason for the small signal in KOI 191 is the large ratio of orbital periods, exceeding 6:1." + Thus. the TTV signal is weakened sieuificautlv.," Thus, the TTV signal is weakened significantly." + If the outermost plauct in IKOI 191 were to have an eccentric orbit then it would eive a periodic TTY signal with a period equal to that of the outer planet as described iu Section | of Agoletal.(2005). (seealsoDorkovitsetal.2003)., If the outermost planet in KOI 191 were to have an eccentric orbit then it would give a periodic TTV signal with a period equal to that of the outer planet as described in Section 4 of \citet{agol2005} \citep[see also][]{bork2003}. +. For KOI 209. the expected TTV sigual for the inner planet shows au abrupt cutoff aud is expected to be larger than a few hundred seconds.," For KOI 209, the expected TTV signal for the inner planet shows an abrupt cutoff and is expected to be larger than a few hundred seconds." + This is because WOT 209.02 has the lougest period of all of the iuner plaucts., This is because KOI 209.02 has the longest period of all of the inner planets. + Caven the time bascline of the extended mission. planets with periods of a few teus of davs will likely prove to be among the most interesting for TTV studies as they siunultaneouslv/ have ongcr periods (the TTY signal is linear in the period) aud will have a sufücieut muuber of transits for a complete analysis.," Given the time baseline of the extended mission, planets with periods of a few tens of days will likely prove to be among the most interesting for TTV studies as they simultaneously have longer periods (the TTV signal is linear in the period) and will have a sufficient number of transits for a complete analysis." + The proximity of IKOI 877 to the 2:1 MMRB lucicates tha this system is likely to have very laree variations., The proximity of KOI 877 to the 2:1 MMR indicates that this system is likely to have very large variations. + However. a steep drop in the expected signal occurs when the orbits are nearly circular.," However, a steep drop in the expected signal occurs when the orbits are nearly circular." + Figure ΤΕ shows an expanded view of the T'TV signal for the inner planet in NOT STI as a function of the iuuer and outer plauct eccentrieities., Figure \ref{877contours} shows an expanded view of the TTV signal for the inner planet in KOI 877 as a function of the inner and outer planet eccentricities. +" From this figure one cau see that. while the zero ecceutricitv case exlibits a relatively αιμα TTV signal. eccentricities nich larger than 0,01 cause the signal to increase bevoud an SNR of unity near LO? seconds (~15 1uuutes)."," From this figure one can see that, while the zero eccentricity case exhibits a relatively small TTV signal, eccentricities much larger than 0.01 cause the signal to increase beyond an SNR of unity near $10^3$ seconds $\sim 15$ minutes)." + Should the TTY sienal © this size or smaller. it should stringently constrain the eccentricities of both planets in the abseuce of auv other data.," Should the TTV signal be this size or smaller, it should stringently constrain the eccentricities of both planets in the absence of any other data." + We note that all of these results for the expected TTV signal have significant depeudence on the eccentrieities of the planets., We note that all of these results for the expected TTV signal have significant dependence on the eccentricities of the planets. + Oue COlisequence of this fact is that. if a large fraction of these or other imultiple svstenms do not show a TTV sienal. then low eccentricity orbits are much more ΟΠΛΟ in unilti-planet systems than iu sinele planet svstenis.," One consequence of this fact is that, if a large fraction of these or other multiple systems do not show a TTV signal, then low eccentricity orbits are much more common in multi-planet systems than in single planet systems." + The threc-plaue system of NOT 152 porteuds the exciting ar challenging studies of svstenis where there are more than two planets aud where multiple planets transit the star., The three-planet system of KOI 152 portends the exciting and challenging studies of systems where there are more than two planets and where multiple planets transit the star. + This svstei js particularly interesting given the relatively close proximity to the [:2:1 uniltibody resonance., This system is particularly interesting given the relatively close proximity to the 4:2:1 multibody resonance. + However. it is unlikely that this svsteu occupies this resonance eiven the estimated orbital periods of the planetsone being estimated from a sinele transit.," However, it is unlikely that this system occupies this resonance given the estimated orbital periods of the planets—one being estimated from a single transit." + For KOT 152. the middle planet is likely to exhibit the largest TTV signalbeing just outside the 2:1 MMRB with an interior plauet aud perhaps just interior to the 2:1 MMRB ofthe exterior One challenge that theee-planet svstenis; such ax WOT 152. pose js the confusion that cal arise from uultiple. competing perturbers in the TTV signa for a particular planet.," For KOI 152, the middle planet is likely to exhibit the largest TTV signal—being just outside the 2:1 MMR with an interior planet and perhaps just interior to the 2:1 MMR of the exterior One challenge that three-planet systems, such as KOI 152, pose is the confusion that can arise from multiple, competing perturbers in the TTV signal for a particular planet." + We present three broad scenarios for consideration in future studies; although other reeimics may exist: 1) nonresonaut/uonresonaut where there is uo 116211 motion conuuensurability between any pair of planets; 2) resonant/nouresonant where oue pair of plauets has a ΠΟΣΤ motion conunuensurability while the other does not. aud 3) resonant/resonaut where anv pair of planets lies near a mca motion conmieusurability.," We present three broad scenarios for consideration in future studies, although other regimes may exist: 1) nonresonant/nonresonant where there is no mean motion commensurability between any pair of planets, 2) resonant/nonresonant where one pair of planets has a mean motion commensurability while the other does not, and 3) resonant/resonant where any pair of planets lies near a mean motion commensurability." + For the first scenario. the TTV signal due to one perturber should be lavegely independent of he TTY signal due to the second perturber.," For the first scenario, the TTV signal due to one perturber should be largely independent of the TTV signal due to the second perturber." + The effect from both perturbers will be of order the yerturber to stellar ass ratio. and therefore nav be comparable.," The effect from both perturbers will be of order the perturber to stellar mass ratio, and therefore may be comparable." + But their coutributions will contribute linearly. to the overall= signal aud he periodicities in the TTY signal due to one erturber will be independent of the periodicities induced by the other., But their contributions will contribute linearly to the overall signal and the periodicities in the TTV signal due to one perturber will be independent of the periodicities induced by the other. + In other words. a Fourier raustorm of the TTV sienal would likely show wo sets of independent reals (seeSteffen2006) hat can be distinenishedOo provided the data have," In other words, a Fourier transform of the TTV signal would likely show two sets of independent peaks \citep[see][]{stef2006} that can be distinguished provided the data have" +using the SFR derived from the SED fitting we notice that NE and C2 are not strongly deviant anymore compared to the other regions.,using the SFR derived from the SED fitting we notice that NE and C2 are not strongly deviant anymore compared to the other regions. +" This shows that great caution must be employed to estimate the SFR as it can influence the results significantly, especially in the case of interacting systems in which the actual SFR can vary rapidly."," This shows that great caution must be employed to estimate the SFR as it can influence the results significantly, especially in the case of interacting systems in which the actual SFR can vary rapidly." +" As mentioned earlier, ? found that starburst galaxies follow a different Schmidt-Kennicutt law than more quiescent galaxies."," As mentioned earlier, \cite{daddi2010a} found that starburst galaxies follow a different Schmidt–Kennicutt law than more quiescent galaxies." + The interaction in Arp 158 increases the turbulence in the system., The interaction in Arp 158 increases the turbulence in the system. + The question is whether different regions in the system also follow different relations., The question is whether different regions in the system also follow different relations. +" In Fig. 6,,"," In Fig. \ref{fig:KS-plot-D10}," + we compare the regions in Arp 158 with the relations found by ?.., we compare the regions in Arp 158 with the relations found by \cite{daddi2010a}. +" We see that similarly to what ? found, we are seeing 2 different regimes of star formation in Arp 158, provided the SFR estimator is accurate."," We see that similarly to what \cite{daddi2010a} found, we are seeing 2 different regimes of star formation in Arp 158, provided the SFR estimator is accurate." +" The first one regroups all regions, except for NE, which are well fitted by a power law with a slope of 1.42."," The first one regroups all regions, except for NE, which are well fitted by a power law with a slope of 1.42." +" Conversely NE presents a much higher SFR surface density for a similar surface density, with an offset which is qualitatively similar to the one found by ? for starburst galaxies."," Conversely NE presents a much higher SFR surface density for a similar surface density, with an offset which is qualitatively similar to the one found by \cite{daddi2010a} for starburst galaxies." +" The offset is slightly larger in the case of ?,, probably as the objects they studied are more extreme than ours."," The offset is slightly larger in the case of \cite{daddi2010a}, probably as the objects they studied are more extreme than ours." + Another important point is that contrary to ? we find this offset while keeping the Xco conversion factor constant., Another important point is that contrary to \cite{daddi2010a} we find this offset while keeping the $_\mathrm{CO}$ conversion factor constant. + Using a smaller conversion factor similar to that used for LIRGs and ULIRGs would only increase the discrepancy., Using a smaller conversion factor similar to that used for LIRGs and ULIRGs would only increase the discrepancy. + What really sets NE apart is not the gas surface density but the high SFR surface density., What really sets NE apart is not the gas surface density but the high SFR surface density. + A possible explanation is that this region is not simply an inflow-driven starburst but that the increased turbulence and a fragmentation into dense clouds strongly increase the SFR surface density for the same gas surface density (???)..," A possible explanation is that this region is not simply an inflow–driven starburst but that the increased turbulence and a fragmentation into dense clouds strongly increase the SFR surface density for the same gas surface density \citep{teyssier2010a,bournaud2010b,bournaud2011a}." + An observational signature of this would be an an excess of the dense gas fraction as observed by ?.., An observational signature of this would be an an excess of the dense gas fraction as observed by \cite{juneau2009a}. +" In order to determine whether the dense gas fraction is higher, some HCN observations are required."," In order to determine whether the dense gas fraction is higher, some HCN observations are required." +" We stress that the use of a lower Xco factor, as is used for ULIRGs for instance, for the NE region would only exacerbate the discrepancy."," We stress that the use of a lower $_\mathrm{CO}$ factor, as is used for ULIRGs for instance, for the NE region would only exacerbate the discrepancy." + The presence of these 2 modes seen in a resolved way in an interacting system shows that its origin does not depend on the global mass or size of the system but that it is rather linked to the physics of the ISM at scales no larger than 1 kpc., The presence of these 2 modes seen in a resolved way in an interacting system shows that its origin does not depend on the global mass or size of the system but that it is rather linked to the physics of the ISM at scales no larger than 1 kpc. +" Indeed, this scale corresponds to the largest gravitational instabilities in the ISM."," Indeed, this scale corresponds to the largest gravitational instabilities in the ISM." + The Jeans length which is of the order of 100—200 pc in nearby spirals increases up to 500--1000 pc in mergers because of higher densities and velocity dispersions., The Jeans length which is of the order of 100–200 pc in nearby spirals increases up to 500–1000 pc in mergers because of higher densities and velocity dispersions. + In addition this scale also corresponds to the injection scale of turbulence in the ISM (??)..," In addition this scale also corresponds to the injection scale of turbulence in the ISM \citep{elmegreen2003a,bournaud2010a}." + In this paper we have studied how properties of star-forming regions vary across an interacting system., In this paper we have studied how properties of star–forming regions vary across an interacting system. +" To do so we have combined an extensive set of archival and proprietary data tracing the molecular gas (CO), the atomic gas (HD, star formation (FUV and 24 ym), and the stellar populations."," To do so we have combined an extensive set of archival and proprietary data tracing the molecular gas (CO), the atomic gas (HI), star formation (FUV and 24 $\mu$ m), and the stellar populations." +" The interacting system shows a complex morphology, the disks of the 2 colliding galaxies having already interpenetrated."," The interacting system shows a complex morphology, the disks of the 2 colliding galaxies having already interpenetrated." + To ascertain the exact nature of the different regions in the interacting system we have also obtained optical spectra., To ascertain the exact nature of the different regions in the interacting system we have also obtained optical spectra. +" In particular we have obtained a firm identification of the nuclei of the merging galaxies, which was still under debate."," In particular we have obtained a firm identification of the nuclei of the merging galaxies, which was still under debate." +" One, to the East, exhibits a starburst, the other one hosts an AGN."," One, to the East, exhibits a starburst, the other one hosts an AGN." +" A third nucleus, to the West, turns out to be a foreground star."," A third nucleus, to the West, turns out to be a foreground star." + A brief description of the regions of interest in Arp 158 is provided hereafter., A brief description of the regions of interest in Arp 158 is provided hereafter. +ONeMg-He bbs at the present.,ONeMg+He DDs at the present. + During the same period. there are no new-born CO+CO DDs.," During the same period, there are no new-born CO+CO DDs." + reftigdol showsthalthecontribulionofstar formaliontolhepresentnumbers(toppancls)andtotalimergedniumbers(boltompancis)ofalliypesof DL , \\ref{fig_cdot} shows that the of star formation to the present numbers (top panels) and total merged numbers (bottom panels) of all types of DDs decreases monotonically as a function of $t_{\rm sf}$. +This is because most DDs from all epochs survive to the present time due to their wide orbital separations., This is because most DDs from all epochs survive to the present time due to their wide orbital separations. + The number of He+He DDs from each epoch decreases with /.; more sharply than for CO+CO DDs. although early star formation provides more He+He DDs than CO+CO DDs.," The number of He+He DDs from each epoch decreases with $t_{\rm sf}$ more sharply than for CO+CO DDs, although early star formation provides more He+He DDs than CO+CO DDs." + A similar situation arises for CO+He and ONeMeg+He- DDs., A similar situation arises for CO+He and ONeMg+He DDs. + This. result .is consistent. with. stellar evolution. and the assumed7 SF prate., This result is consistent with stellar evolution and the assumed SF rate. +" The significanceEN of. computing. the present number of. DDs using; 9,98refeq, umberl.; whichrepresentsthesumof DDsarisingfromdif ferentstar9.975 formalioncpochs. islhalitdemonslralesthelinkbebwecnthes FhistókyMFElyegalaa (9 ΑΣΕ ΠΛ ΜΜion)andt hec, deyDDs.whicheanbededuced from. forcrample. theirgravitationalwavesigneal,"," The significance of computing the present number of DDs using \\ref{eq_number1}, which represents the sum of DDs arising from different star-formation epochs, is that it demonstrates the link between the SF history of the galaxy (or, at least, the thin disc in the present investigation) and the distribution of the properties of present-day DDs, which can be deduced from, for example, their gravitational wave signal." + rnmishowsthevariationof ς and M nmberly)ofdif ferenttgpesof v. ο Μα.," \\ref{fig_birnum} shows the variation of $\nu$, $\zeta$ and $n_{\rm dd}$ \\ref{eq_number1}) ) of different types of DD with age $t_{\rm disc}$." + The individual properties of the current rad DDs will be used to calculate*aleul the gravitationalepavitatii wave“ave signalspon ?(2).., The individual properties of the current $n_{\rm dd}$ DDs will be used to calculate the gravitational wave signal \citep{Yu11}. + The use of a realistic dise model is important in order to describe the distance distribution of white dwarf binary systems from the Sun., The use of a realistic disc model is important in order to describe the distance distribution of white dwarf binary systems from the Sun. + ? proposed a double exponential distribution., \citet{Sackett97} proposed a double exponential distribution. + ? derived three functions for the star density distribution in their model of a thin disc plus thick disc (exponential | exponential. hyperbolic secant | exponential. and squared hyperbolic secant || exponential. respectively) from fits to deep star counts carried out in the Calar Alto Deep Imaging Survey.," \citet{Phleps00} derived three functions for the star density distribution in their model of a thin disc plus thick disc (exponential $+$ exponential, hyperbolic secant $+$ exponential, and squared hyperbolic secant $+$ exponential, respectively) from fits to deep star counts carried out in the Calar Alto Deep Imaging Survey." + We here model the thin dise in the Galaxy using a squared hyperbolic secant plus exponential distribution expressed as: where // and z are the natural eylindrical coordinates of the axisymmetric dise. j;=2.5 Kkpe is the scale length of the disc. and fy.=0.352 kkpe is the scale height of the thin dise.," We here model the thin disc in the Galaxy using a squared hyperbolic secant plus exponential distribution expressed as: where $R$ and $z$ are the natural cylindrical coordinates of the axisymmetric disc, $h_{R}=2.5$ kpc is the scale length of the disc, and $h_{z}=0.352$ kpc is the scale height of the thin disc." +" AZ, is the mass of the thin disc. which is determined by the star formation rate."," $M_{\rm tn}$ is the mass of the thin disc, which is determined by the star formation rate." + We adopt the position of the Sun to be Rau=S. KKkpe. zou = 16.5 Προ (2)..," We adopt the position of the Sun to be $R_{\rm sun} = 8.5$ kpc, $z_{\rm sun}$ = $16.5$ pc \citep{Freudenreich98}." + We neglect the age and mass dependence of the scale height., We neglect the age and mass dependence of the scale height. + This thin-dise model is consistent with the model of ? and ?.. and also in agreement with Hipparcos results and the observed rotation curve.," This thin-disc model is consistent with the model of \citet{Klypin02} and \citet{Robin03}, and also in agreement with Hipparcos results and the observed rotation curve." +" From the SF rate. the total mass of stars in the thin dise at age 10 Gyr is Afi,z 107AL.."," From the SF rate, the total mass of stars in the thin disc at age 10 Gyr is $M_{\rm tn}\approx$ $\times10^{10}~M_{\odot}$." + Combining the thin disc model and the mass of stars in the thin disc. the stellar density in the solar neighbourhood is 6.27107M.pe? for the thin disc. These values are consistent with the Hipparcos result. (7.6d1.5)/107?M.pe 72(2) and the dynamical structure of the thin disc (2?)..," Combining the thin disc model and the mass of stars in the thin disc, the stellar density in the solar neighbourhood is $6.27\times10^{-2}{\rm M_{\odot}pc^{-3}}$ for the thin disc, These values are consistent with the Hipparcos result, $7.6\pm1.5)\times10^{-2}~ +{\rm M_{\odot}pc^{-3}}$ \citep{Creze98} and the dynamical structure of the thin disc \citep{Klypin02,Robin03}." + The local density of DDs in the model is L98«1074 pe., The local density of DDs in the model is $1.98\times10^{-4}$ $\rm pc^{-3}$. + The presence of unevolved DDs in facrilicallime when the rogeniorindicalestheeristenceof puxeceitgted cthe, The presence of unevolved DDs in \\ref{fig_progenitor} indicates the existence of a critical time when the first DD of each type was just born. +" reftig,πας| CO.CO|He andtLe UeDDsrespeelively. thetimesare25 Alyy. DOAL gr. SOOAL gr. and650AL yr."," For ONeMg+X, CO+CO, CO+He, and DDs respectively, the times are 25 Myr, 50 Myr, 560 Myr, and 650 Myr." + Figipillustralesthecontribulionofdi fferentepochsof star formeationtoli dagdistribulionoflolalmassandorbilalperiodsofthe DDs, \\ref{fig_mp} illustrates the contribution of different epochs of star formation to the present-day distribution of total mass and orbital periods of the DDs. + Meheredistinguis 6.026 ⋅-2.5$ ) is sensitive to star formation between 8 and 9.95 Gyr after the thin disc formed, which means that these DDs are most likely to be young." + Their MS+MS progenitors formed between 50 Myr and 2000 Myr ago., Their MS+MS progenitors formed between 50 Myr and 2000 Myr ago. +" The total stellar mass formed during the time represented by each panel of reftig,,, p/83.62. 0.85. 0.38. 0.15.0.19. @7d0.0087- 102 AL. ta to T."," The total stellar mass formed during the time represented by each panel of \\ref{fig_mp} is $3.62$, $0.85$, $0.38$, $0.15$, $0.19$, and $0.0087\times10^{10}$ $M_{\odot}$ (a to f)." + The current number of DDs in the thin disc derived from euch star formation in the figure is given in table 2.., The current number of DDs in the thin disc derived from each star formation epoch in the figure is given in table \ref{tab_DDsf}. + These numbers indicate that. for epochcurrent He+He and CO+He DDs. a large number and 87.856)) have ages greater than + Gyr. while only and of CO+CO and ONeMg+X have ages in the same range.," These numbers indicate that, for current He+He and CO+He DDs, a large number and ) have ages greater than 4 Gyr, while only and of CO+CO and ONeMg+X have ages in the same range." + A significant number and 6.7569) of CO+CO and ONeMg+X DDs have been produced by the last | Gyr of star formation., A significant number and ) of CO+CO and ONeMg+X DDs have been produced by the last 1 Gyr of star formation. + The number of He+tHe and CO+He DDs from this period is negligible., The number of He+He and CO+He DDs from this period is negligible. + However. ?. show that DDs with ages less than 2 Gyr would contribute substantially to the amplitude of the gravitational wave signal in several frequeney bands.," However, \citet{Yu11} show that DDs with ages less than 2 Gyr would contribute substantially to the amplitude of the gravitational wave signal in several frequency bands." + According to the classification of the star formation stages and the eritical time for the birth of DDs. we can see from plhaltheALS| AlSprogenitorsofCO | CODDsare formedbe forety=9.95 Gyr.," According to the classification of the star formation stages and the critical time for the birth of DDs, we can see from \\ref{fig_mp} that the MS+MS progenitors of CO+CO DDs are formed before $t_{\rm sf}\approx9.95$ Gyr." + The youngest CO+He DD has an age of about 560 Myr. but the majority of their MS progenitors formed 7600 Myr ago.," The youngest CO+He DD has an age of about 560 Myr, but the majority of their MS progenitors formed $>$ 600 Myr ago." + These results are consistent with stellar evolution calculations., These results are consistent with stellar evolution calculations. + Note that in the stellar evolution model a fraction of ONeM;z white dwarfs become neutron stars and. stellar-mass black holes due to accretion-induced collapse., Note that in the stellar evolution model a fraction of ONeMg white dwarfs become neutron stars and stellar-mass black holes due to accretion-induced collapse. + These do not form type Ia supernovae and are not otherwise considered in our results., These do not form type Ia supernovae and are not otherwise considered in our results. + with a. quasi- declining SF rate., We have simulated the present DD population with a quasi-exponential declining SF rate. + In presentorder to see the /IDDChiNpuldtidninfluence on the, In order to see the influence on the +This corresponds (to a mass AMgu70.52M... πο that q~0.41.,"This corresponds to a mass $M_{Ab} \sim 0.52$, so that $q \sim 0.41$." + We have presented the first evidence of total eclipses (including eclipses of the faint secondarv star in { band) in the svstem $986 using an extensive series of photometric observations., We have presented the first evidence of total eclipses (including eclipses of the faint secondary star in $I$ band) in the system S986 using an extensive series of photometric observations. + We have taken high-resolution spectroscopy. of the system and have identified a (hind star that contributes to (he light of the svstem., We have taken high-resolution spectroscopy of the system and have identified a third star that contributes to the light of the system. + The (third star appears to be a cluster meniber. but may or may not be physically associated will the eclipsing binary.," The third star appears to be a cluster member, but may or may not be physically associated with the eclipsing binary." + The results ol our analvsis are given in Table 5.., The results of our analysis are given in Table \ref{props}. + The detailed analvsis of our spectroscopy and photometry for the $986 indicates that the primary star in the eclipsing binary (component Aa) is a star that is slightlv (bul significantlv) hotter than the Curnoff of the cluster., The detailed analysis of our spectroscopy and photometry for the S986 indicates that the primary star in the eclipsing binary (component Aa) is a star that is slightly (but significantly) hotter than the turnoff of the cluster. + Two stellar explanations for this exist., Two stellar explanations for this exist. + One possibility is that the primary is a normal main sequence star (hat is in a relatively short lived phase of its evolution., One possibility is that the primary is a normal main sequence star that is in a relatively short lived phase of its evolution. + The gap in the CMD of M67 with 12.85=0.616+ 0.046).","), the hypothesis of a uniform distribution of the objects in the Universe is rejected at a confidence level of $\sim 98\%$ $ = 0.616 \pm 0.046$ )." + Assuming a pure luminosity evolution model with an evolutionary form ος(1+zY€ we finda best fit parameter of C~2.7 with an associated confidence interval of 1.9-3.0., Assuming a pure luminosity evolution model with an evolutionary form $\propto (1+z)^C$ we find a best fit parameter of $\simeq$ 2.7 with an associated confidence interval of 1.9-3.0. +" This value of cosmological evolution is consistent, within the errors, both with the results obtained in the soft (E<3 keV) energy band using the Extended Medium Sensitivity Survey (C=2.56+0.17; Maccacaroet 1991,, DellaCecaetal. 1992)) and the EMSS+Rosat AGN samples (C=2.6+0.1; Pageetal. 1997)) and with the results in the 2-10 keV energy range reported in Uedaetal.(2003) (C= 2.70*057) and LaFrancaetal.(2005) (C= 3.2205)."," This value of cosmological evolution is consistent, within the errors, both with the results obtained in the soft $\ls 3$ keV) energy band using the Extended Medium Sensitivity Survey $C=2.56\pm 0.17$; \citealt{maccacaro1991}, , \citealt{dellaceca1992}) ) and the EMSS+Rosat AGN samples $C=2.6\pm 0.1$; \citealt{page1997}) ) and with the results in the 2-10 keV energy range reported in \cite{ueda2003} $C=2.70^{+0.17}_{-0.25}$ ) and \cite{lafranca2005} $C=3.22^{+0.13}_{-0.26}$ )." +" However it is now well established that a Luminosity Dependent Density Evolution (LDDE) model provides a better description of the evolutionary properties of AGN, both in the X-ray energy range (Hasingeretal.2005;; Uedaetal.2003;; LaFrancaetal.2005;; Silvermanetal.2007)) and in the optical domain (Bongiornoetal.2007))."," However it is now well established that a Luminosity Dependent Density Evolution (LDDE) model provides a better description of the evolutionary properties of AGN, both in the X-ray energy range \citealt{hasinger2005}; \citealt{ueda2003}; \citealt{lafranca2005}; \citealt{silverman2007}) ) and in the optical domain \citealt{bongiorno2007}) )." +" To test this evolutionary behavior, we assume here an LDDE model with the parametrization as introduced by Uedaetal.(2003), where Ζς corresponds to the redshift where the direction of the evolution changes sign."," To test this evolutionary behavior, we assume here an LDDE model with the parametrization as introduced by \cite{ueda2003}, where $_{\rm c}$ corresponds to the redshift where the direction of the evolution changes sign." +" It is worth noting that z, is a function of the intrinsic luminosity of the object; if we assume the best fit parameters of p2=-1.15; z;=2.49; a=0.20; Log L,=45.80 (adapted to Ho=65) as reported in LaFrancaetal.(2005),, then z, is 0.7,1.1,1.7 for AGN with Lx 10%,10*,, respectively."," It is worth noting that $_{\rm c}$ is a function of the intrinsic luminosity of the object; if we assume the best fit parameters of p2=-1.15; $_{\rm c}^*=2.49$; $\alpha$ =0.20; Log $_a$ =45.80 (adapted to $_0$ =65) as reported in \cite{lafranca2005}, then $_{\rm c}$ is $\sim 0.7,1.1,1.7$ for AGN with $_{\rm X}$ $\sim 10^{43}, 10^{44}, 10^{45}$, respectively." +" Given the coverage in the luminosity-redshift plane of the HBSS unabsorbed AGN sample, for each luminosity the objects are below z;, implying that we are unable to derive p2, z;, a and LogL,."," Given the coverage in the luminosity-redshift plane of the HBSS unabsorbed AGN sample, for each luminosity the objects are below $_{\rm c}$, implying that we are unable to derive p2, $_{\rm c}^*$, $\alpha$ and $_a$ ." +" For this reasons we have fixed them from LaFrancaetal.(2005) and we have used the V,/V, test to constrain pl.", For this reasons we have fixed them from \cite{lafranca2005} and we have used the $V_e/V_a$ test to constrain p1. + We obtain a best fit pl=6.5 with an associated confidence interval of 3.5 - 10.0., We obtain a best fit p1=6.5 with an associated confidence interval of 3.5 - 10.0. +" The distribution of the derived V,/V, values is consistent with being uniformly distributed between 0 and 1 according to a KS test (KS probability 9596).", The distribution of the derived $V_e/V_a$ values is consistent with being uniformly distributed between 0 and 1 according to a KS test (KS probability $\sim 95\%$ ). +" We have also checked that, given the coverage of the luminosity-redshift plane of the HBSS AGN sample, the best fit pl is virtually insensitive to the other parameters of the model; ie. pl does not change by varying all the other parameters within their lo range as derived from LaFrancaetal.(2005)."," We have also checked that, given the coverage of the luminosity-redshift plane of the HBSS AGN sample, the best fit p1 is virtually insensitive to the other parameters of the model; i.e., p1 does not change by varying all the other parameters within their $1\sigma$ range as derived from \cite{lafranca2005}." +". The derived best fit value for pl is consistent, within the errors, with that reported in LaFrancaetal.(2005) (ρ1--4.62+ 0.26) and is in very good agreement with those recently obtained, in the optical domain, by Bongiornoet(2007) using a sample of 130 broad line AGN with redshift up to z=5 from the VIMOS-VLT Deep Survey (p1=6.54) and from Hopkinsetal.(2007) using a large data set of AGN selected in the Mid-IR, optical, soft X-ray and hard X-ray (p1=5.95+0.23)."," The derived best fit value for p1 is consistent, within the errors, with that reported in \cite{lafranca2005} $4.62\pm 0.26$ ) and is in very good agreement with those recently obtained, in the optical domain, by \cite{bongiorno2007} using a sample of 130 broad line AGN with redshift up to z=5 from the VIMOS-VLT Deep Survey (p1=6.54) and from \cite{hopkins2007} using a large data set of AGN selected in the Mid-IR, optical, soft X-ray and hard X-ray $5.95\pm 0.23$ )." + Because of their number statistics (22 objects in total) and their distribution in the Ly—z plane the cosmological evolution is unconstrained for the absorbed AGN sample (note that the absorbed AGN are sampled only up to z~0.8)., Because of their number statistics (22 objects in total) and their distribution in the $L_X-z$ plane the cosmological evolution is unconstrained for the absorbed AGN sample (note that the absorbed AGN are sampled only up to $\sim 0.8$). +" Therefore in the following, and in line with the Unification Scheme of AGN, we will make the assumptionthat this class of sources evolve withcosmic time (and within the reshift range sampled at the HBSS flux limit) in a similar way as the unabsorbed ones."," Therefore in the following, and in line with the Unification Scheme of AGN, we will make the assumptionthat this class of sources evolve withcosmic time (and within the reshift range sampled at the HBSS flux limit) in a similar way as the unabsorbed ones." +We are particularly eratelul to E.,We are particularly grateful to E. +"data (Φ”=0.00016+0.0004Mpc,0.3Mgyr!,a=—1.51 40.08) gives psrr=(25E1.7)x107?Meyt! Μρς ","data $\Phi^* = 0.00016\pm0.0004\; \mathrm{Mpc}^{-3},\; \psi^*=9.2\pm0.3\;\mathrm{M}_{\sun} \; \mathrm{yr}^{-1}, \;\alpha=-1.51\pm0.08$ ) gives $\rho_\mathrm{SFR} = (25\pm1.7) \times 10^{-3} \;\mathrm{M}_{\sun}\; \mathrm{yr}^{-1}\; \mathrm{Mpc}^{-3}$ ." +This is in good agreementὃ. with most recent derivations of this result: see Table 1 for a compilation of recent results., This is in good agreement with most recent derivations of this result: see Table \ref{tab1} for a compilation of recent results. +" There is a relatively large spread in the derived values of the SFR volume density - greater than a factor of two, beyond the errors quoted on the individual measurements."," There is a relatively large spread in the derived values of the SFR volume density - greater than a factor of two, beyond the errors quoted on the individual measurements." +" This is discussed briefly by ? (who derive their own 1.4 GHz-based value of (21+5)x107?Meyr! 5), who attribute the discrepancy to a systematic underestimationΜρςε of the extinction using the Balmer decrement in some emission line-based studies."," This is discussed briefly by \cite{2002MNRAS.330..621S} (who derive their own 1.4 GHz-based value of $(21\pm5) \times 10^{-3}\;\mathrm{M}_{\sun}\; \mathrm{yr}^{-1}\; \mathrm{Mpc}^{-3}$ ), who attribute the discrepancy to a systematic underestimation of the extinction using the Balmer decrement in some emission line-based studies." +" The total value of psrr can also be decomposed into ‘UV’ and ‘IR’ components, by integrating the value of i?Φ(Φ) derived from each component individually."," The total value of $\rho_\mathrm{SFR}$ can also be decomposed into `UV' and `IR' components, by integrating the value of $\psi \,\Phi(\psi) $ derived from each component individually." +" Doing so leads to values of psrr(IR)=0.011Meyr!Mpc?, and psreR(UV)=0.012MeyrMpc?."," Doing so leads to values of $\rho_\mathrm{SFR}(\mathrm{IR}) = 0.011\; \mathrm{M}_{\sun}\; \mathrm{yr}^{-1}\; \mathrm{Mpc}^{-3}$, and $\rho_\mathrm{SFR}(\mathrm{UV}) = 0.012\; \mathrm{M}_{\sun}\; \mathrm{yr}^{-1}\; \mathrm{Mpc}^{-3}$." + The LVL contribution is 0.0007Mayr!Mpc?., The LVL contribution is $0.0007\; \mathrm{M}_{\sun}\; \mathrm{yr}^{-1}\; \mathrm{Mpc}^{-3}$. +" This is4796,50%,, of the total for the IR, UV, and LVL components respectively."," This is, of the total for the IR, UV, and LVL components respectively." +" This result - that about half of the energy from the total cosmic star formation budget is re-processed by dust - is well known, and is in line with previous studies."," This result - that about half of the energy from the total cosmic star formation budget is re-processed by dust - is well known, and is in line with previous studies." +" ? found that of their derived total SFR. volume density (19x10?MoyrΜρο ?), was from dust-reprocessed For !consistency (and because our statistical AGN removal involves some uncertainty), we have checked the value of psrn calculated from the sample the statistical correction for AGN contamination (as per 833.1)."," \cite{Takeuchi:2005aa} found that of their derived total SFR volume density $19 \times 10^{-3} \; \mathrm{M}_{\sun}\; \mathrm{yr}^{-1}\; \mathrm{Mpc}^{-3}$ ), was from dust-reprocessed For consistency (and because our statistical AGN removal involves some uncertainty), we have checked the value of $\rho_\mathrm{SFR}$ calculated from the sample the statistical correction for AGN contamination (as per 3.1)." +" As the correction is only significant at the upper end (beyond φ”), the value only changes slightly: without any AGN correction applied, we calculate psrr=(26+2.2)x107?Meyt! Μρς "," As the correction is only significant at the upper end (beyond $\psi^*$ ), the value only changes slightly: without any AGN correction applied, we calculate $\rho_\mathrm{SFR} = (26\pm2.2) \times 10^{-3}\;\mathrm{M}_{\sun}\; \mathrm{yr}^{-1}\; \mathrm{Mpc}^{-3}$ ." +We may also ὃ.compute the fraction of the local cosmic star formation rate density occurring in starburst environments., We may also compute the fraction of the local cosmic star formation rate density occurring in starburst environments. +" For the purposes of such an analysis, we define a starburst as a system forming stars at 210 yr|."," For the purposes of such an analysis, we define a starburst as a system forming stars at $\geq$ 10 $_{\sun}\; \mathrm{yr}^{-1}$ ." +" Using the star formation rate density distribution,Mo we can thus integrate from 10 Meyr! to infinity: For our data, this value is 0.0049+0.00039 Moyr!Mpc?, or of the total star formation rate volume density; by our (admittedly somewhat crude) definition, one fifth of the starformation in the local Universe is provided by starbursts."," Using the star formation rate density distribution, we can thus integrate from 10 $_{\sun}\; \mathrm{yr}^{-1}$ to infinity: For our data, this value is $0.0049\pm 0.00039$ $_{\sun}\; \mathrm{yr}^{-1}\; \mathrm{Mpc}^{-3}$, or of the total star formation rate volume density; by our (admittedly somewhat crude) definition, one fifth of the starformation in the local Universe is provided by starbursts." + This is consistent with the values found by ? using specific star formation rates from SDSS., This is consistent with the values found by \cite{2004MNRAS.351.1151B} using specific star formation rates from SDSS. +" Interestingly, ? also find that of star formation in the dwarf galaxy population is concentrated in high Ho equivalent width systems."," Interestingly, \cite{2009ApJ...692.1305L} also find that of star formation in the dwarf galaxy population is concentrated in high $\alpha$ equivalent width systems." +" 'There are many interesting values that can be derived from the distribution shown in Fig. 5,"," There are many interesting values that can be derived from the distribution shown in Fig. \ref{fig:sfr_den}," +" including the starburst fraction (discussed above), the ‘dividing’ SFR at which of the star formation is happening both above and below, and so on."," including the starburst fraction (discussed above), the `dividing' SFR at which of the star formation is happening both above and below, and so on." +" Rather than providng list of values for various integration limits, it is more enlighteninga to consider the behaviour of the cumulative fraction of star formation rate volume density, which is shown in Fig. 6.."," Rather than providng a list of values for various integration limits, it is more enlightening to consider the behaviour of the cumulative fraction of star formation rate volume density, which is shown in Fig. \ref{fig:cum_sfr_den}." +" This shows SFR, plotted against the fraction of the total star formation volume density coming from SFRs than that SFR."," This shows SFR, plotted against the fraction of the total star formation volume density coming from SFRs than that SFR." +" The data show a power-law increase in star formation rate volume density fraction, over 5 orders of magnitude until the truncation at ~20Mc,yr’."," The data show a power-law increase in star formation rate volume density fraction, over 5 orders of magnitude until the truncation at $\sim20 \;\mathrm{M}_{\sun}\; \mathrm{yr}^{-1}$." +" From this it can be seen that the *5076"" divide occurs at ~3M;yr!, about the SFR of the Milky Way (e.g. ?))."," From this it can be seen that the ' divide occurs at $\sim 3\; \mathrm{M}_{\sun}\; \mathrm{yr}^{-1}$, about the SFR of the Milky Way (e.g. \citealt{2006A&A...459..113M}) )." + It is also interesting to consider the contribution to the total star formation rate volume density from LIRGs and ULIRGs., It is also interesting to consider the contribution to the total star formation rate volume density from LIRGs and ULIRGs. +" These IR-bright galaxies (defined as havingLm>10!Lc and >1013Lo, respectively) are rare in the local Universe, but become more and more important with lookback time, becoming an increasingly dominant componentof the total star formation rate volume density at higher redshifts (?;; ?;; ?))."," These IR-bright galaxies (defined as having$\mathrm{L}_{\mathrm{IR}} > 10^{11}\;\mathrm{L}_{\sun}$ and $> 10^{12}\;\mathrm{L}_{\sun}$ respectively) are rare in the local Universe, but become more and more important with lookback time, becoming an increasingly dominant componentof the total star formation rate volume density at higher redshifts \citealt{2005ApJ...619L..47S}; ; \citealt{2009A&A...496...57M}; ; \citealt{2010arXiv1008.0859G}) )." + Fig., Fig. + 7 shows the star formation rate distribution, \ref{fig:ULIRG} shows the star formation rate distribution +"Here R, P, and B1» are the neutron star radius, rotation period (in s), and magnetic field (in 1013 G), respectively.","Here $R$ , $P$ , and $B_{12}$ are the neutron star radius, rotation period (in s), and magnetic field (in $10^{12}$ G), respectively." +" Accordingly, yioo=7/100, vaHz is the wave frequency in GHz, and A,=A/10!, where A=me/ngj is the multiplicity of the particle creation near magnetic poles (na; ΩΒ/2ποε is the Goldreich-Julian number density)."," Accordingly, $\gamma_{100} = \gamma/100$, $\nu_{\rm GHz}$ is the wave frequency in GHz, and $\lambda_{4} = \lambda/10^{4}$, where $\lambda = n_{\rm e}/n_{\rm GJ}$ is the multiplicity of the particle creation near magnetic poles $n_{\rm GJ} = \Omega B/2 \pi c e$ is the Goldreich-Julian number density)." +" On the other hand, the transverse extraordinary wave with the refractive index (X-mode) is to propagate freely."," On the other hand, the transverse extraordinary wave with the refractive index (X-mode) is to propagate freely." + As the radius rA is much smaller than the escape radius resc (Cheng Ruderman 1979; Andrianov Beskin 2010) one can consider the effects of refraction and limiting polarization separately., As the radius $r_{\rm A}$ is much smaller than the escape radius $r_{\rm esc}$ (Cheng Ruderman 1979; Andrianov Beskin 2010) one can consider the effects of refraction and limiting polarization separately. +" In particular, this implies that one can consider the propagation of waves in the region rresc as rectilinear."," In particular, this implies that one can consider the propagation of waves in the region $r \sim r_{\rm esc}$ as rectilinear." +" Below for simplicity we assume that both two outgoing modes are generated at the same heights rem (few to tens NS radii), where the magnetic field can be considered as a rotating dipole Here = is the corresponding pulsar rotation phase."," Below for simplicity we assume that both two outgoing modes are generated at the same heights $r_{\rm em}$ (few to tens NS radii), where the magnetic field can be considered as a rotating dipole Here = is the corresponding pulsar rotation phase." +"2),we have we have In the rotating vector model (RVM) the p.a. is determined purely by the projection of magnetic field on the sky’s plane, so it coincides with $;,.",",we have we have In the rotating vector model (RVM) the ${\it p.a.}$ is determined purely by the projection of magnetic field on the sky's plane, so it coincides with $\phi_{m}$." +" The sign of the arctan term is determined by the p.a. measuredcounter-clockwise in the picture plane, as is common in radio astronomy (Everett Weisberg 2001)."," The sign of the arctan term is determined by the ${\it p.a.}$ measured in the picture plane, as is common in radio astronomy (Everett Weisberg 2001)." +" As the aberration angle at the emission point is approximately Qrem/c, i.e., it is much smaller than the angular size of the emission cone 1/», we can easily find the position of the emission point, at which the magnetic field line is along the line of sight."," As the aberration angle at the emission point is approximately $\Omega r_{\rm em}/{c}$, i.e., it is much smaller than the angular size of the emission cone $1/\gamma$, we can easily find the position of the emission point, at which the magnetic field line is along the line of sight." +" This point rem=(rem;Jem;Pem) in the XYZ frame is given by the spherical angles as Note that the impact angle )isthesmallestanglebetweenlineof sightandmagneticmomentm, isgivenby — )."," This point ${\bf r_{\rm em}} = (r_{\rm em},\theta_{\rm em},\phi_{\rm em})$ in the $XYZ$ frame is given by the spherical angles as Note that the impact angle is the smallest angle between line of sight and magnetic moment ${\bf m}$, is given by -." +".Asaresult, thetrajectoryoftheextraordinarywaveintheXY 7 f rameisgiver 'This relation allows us to determine magnetic field and all plasma characteristics along the ray."," As a result, the trajectory of the extraordinary wave in the $XYZ$ frame isgivenby the simple relation This relation allows us to determine magnetic field and all plasma characteristics along the ray." +range covered by the data.,range covered by the data. + In those cases. the tail of the broad gaussian just serves to attenuate the emission of the source at long wavelengths.," In those cases, the tail of the broad gaussian just serves to attenuate the emission of the source at long wavelengths." + Furthermore. we find no clear trend in the central energy or the width of the line.," Furthermore, we find no clear trend in the central energy or the width of the line." + We find that the ratio of the spectra of the last and the first observation is close to a power law., We find that the ratio of the spectra of the last and the first observation is close to a power law. + We therefore fit the data to an empirical model that consists of a blackbody multiplied by a power law ET. all affected by interstellar absorption.," We therefore fit the data to an empirical model that consists of a blackbody multiplied by a power law $E^{\Gamma}$, all affected by interstellar absorption." + While the index of the multiplicative power law is allowed to change between observations. for these fits we constrain the parameters of the blackbody and the interstellar absorption to be the same in all observations.," While the index of the multiplicative power law is allowed to change between observations, for these fits we constrain the parameters of the blackbody and the interstellar absorption to be the same in all observations." + While it is difficult to assign a physical interpretation to this model. it provides an acceptable description of the data in the RGS range ( 10 - 38A.. see Table 2). it has fewer parameters than the gaussian absorption model and. in addition. the index of the power law increases steadily over the course of the observations.," While it is difficult to assign a physical interpretation to this model, it provides an acceptable description of the data in the RGS range ( 10 - 38, see Table 2), it has fewer parameters than the gaussian absorption model and, in addition, the index of the power law increases steadily over the course of the observations." + Starting from the raw data. we first produce a lst of calibrated events.," Starting from the raw data, we first produce a list of calibrated events." + To reduce pile-up. in the next step we select only single events as well as events that are not affected by some of the imperfections (bad columns. hot pixels. ete.)," To reduce pile-up, in the next step we select only single events as well as events that are not affected by some of the imperfections (bad columns, hot pixels, etc.)" + of the CCDs., of the CCDs. + We extract events withina 39 aresee circle centred on the source., We extract events withina $39$ arcsec circle centred on the source. + We barycenter these events using the SAS routine BARYCEN version 1.13.4. and we then separate the events according to their energy in. 3 event lists: the bands that we use are 0.1 to 1.2 keV. O.I to 0.4 keV. and 0.4 to 0.8 keV. respectively.," We barycenter these events using the SAS routine BARYCEN version 1.13.4, and we then separate the events according to their energy in 3 event lists; the bands that we use are 0.1 to 1.2 keV, 0.1 to 0.4 keV, and 0.4 to 0.8 keV, respectively." + For each observation we find the best period in the full band using an epoch folding technique: m all cases we find a period of 8.39] s. consistent with the value previously found for this source by ?..," For each observation we find the best period in the full band using an epoch folding technique; in all cases we find a period of 8.391 s, consistent with the value previously found for this source by \citet{kaplan03}." + We then produce folded light curves in the three bands. and we also compute a folded hardness-ratio light curve from the ratio of the 0.4-0.8 keV and the 0.1—0.4 keV light curves.," We then produce folded light curves in the three bands, and we also compute a folded hardness-ratio light curve from the ratio of the 0.4–0.8 keV and the 0.1–0.4 keV light curves." + In Figure 2 we show the 0.1-1.2 keV and the hardness- light curves., In Figure \ref{pulse} we show the 0.1-1.2 keV and the hardness-ratio light curves. + For each observation we define the phase such that the maximum of the full-band light curve occurs at phase zero: the phase of the hardness ratio light curves is the same as for the full-band light curves., For each observation we define the phase such that the maximum of the full-band light curve occurs at phase zero; the phase of the hardness ratio light curves is the same as for the full-band light curves. + The pulse profile in the 0.1—1.2 keV band. as well as the hardness-ratio pulse profile. change from one observation to the other.," The pulse profile in the 0.1–1.2 keV band, as well as the hardness-ratio pulse profile, change from one observation to the other." + The first panel in Figure 2 shows a sinusoidal fit to the pulse profile during the first observation; the same sine function is overplotted to the full-band pulse profiles obtained from the other observations., The first panel in Figure \ref{pulse} shows a sinusoidal fit to the pulse profile during the first observation; the same sine function is overplotted to the full-band pulse profiles obtained from the other observations. + It is apparent that the pulse profile becomes narrower with time., It is apparent that the pulse profile becomes narrower with time. + At the same time. the hardness-ratio pulse. profile also changes.," At the same time, the hardness-ratio pulse profile also changes." + lr the first observation there is a clear modulation. and the hardness-ratio profile leads the full-band light curve by (0.061+0.017 in phase.," In the first observation there is a clear modulation, and the hardness-ratio profile leads the full-band light curve by $0.064 \pm 0.017$ in phase." + In the following observations the amplitude of the hardness-ratio modulation decreases and the phase difference between the full-band and the hardness-ratio light curves is consistent with zero., In the following observations the amplitude of the hardness-ratio modulation decreases and the phase difference between the full-band and the hardness-ratio light curves is consistent with zero. + Eventually. in the last observation the modulation increases again. but now the hardness-ratio light curve lags the full band-light curve by 0.126+0.010 in phase.," Eventually, in the last observation the modulation increases again, but now the hardness-ratio light curve lags the full band-light curve by $-0.126 \pm 0.010$ in phase." + The ddata of sshow that the spectrum of the source changes on a time scale of years. the first time ever that the X-ray spectrum of an isolated neutron star. other then soft gamma-ray repeaters or anomalous X-ray pulsars. is seen to change.," The data of show that the spectrum of the source changes on a time scale of years, the first time ever that the X-ray spectrum of an isolated neutron star, other then soft gamma-ray repeaters or anomalous X-ray pulsars, is seen to change." + Whereas the changes are most pronounced in the last observation. we think that the actual change ts gradual. as witnessed by a gradual increase in the temperatures derived from the blackbody fits: or by a gradual increase in the index of the powerlaw in the fits with a blackbody multiplied with a power law reftemp and reffluxed)).," Whereas the changes are most pronounced in the last observation, we think that the actual change is gradual, as witnessed by a gradual increase in the temperatures derived from the blackbody fits; or by a gradual increase in the index of the powerlaw in the fits with a blackbody multiplied with a power law \\ref{temp} and \\ref{fluxed}) )." + The spectral changes are accompanied by an energy-dependent change in the pulse shape: in particular the pulse phase where the spectrum is hardest has moved with respect to the phase of maximum flux refpulse))., The spectral changes are accompanied by an energy-dependent change in the pulse shape; in particular the pulse phase where the spectrum is hardest has moved with respect to the phase of maximum flux \\ref{pulse}) ). + The phase aangle) dependent spectrum of single neutron stars is currently not explained., The phase angle) dependent spectrum of single neutron stars is currently not explained. + The broad absorption features have been interpreted as a proton-cyclotron absorption feature (?).., The broad absorption features have been interpreted as a proton-cyclotron absorption feature \citep{haberl03a}. . + In pulsars with a strong field (probably stronger than the limit for JO720.4-3125)) the absorption feature, In pulsars with a strong field (probably stronger than the limit for ) the absorption feature +"a dispersion of 0.1 magnitudes and the observational errors as 0.23 magnitudes, which are approximately the averages of the observational errors of the presently available SNe Ia (???)..","a dispersion of $0.1$ magnitudes and the observational errors as $0.23$ magnitudes, which are approximately the averages of the observational errors of the presently available SNe Ia \citep{Riess:2006fw, WoodVasey:2007jb, Davis:2007na}." + The total error in our generated distance moduli for SN Ia is therefore V0.1?+0.23? magnitudes., The total error in our generated distance moduli for SN Ia is therefore $\sqrt{0.1^2+0.23^2}$ magnitudes. + We assume an uniform distribution for SNe Ia along the redshifts., We assume an uniform distribution for SNe Ia along the redshifts. +" For GRBs, instead of generating mock data about the five luminosity relations (see ?)), we directly generate distance moduli like we do for SNe Ia for simplicity."," For GRBs, instead of generating mock data about the five luminosity relations (see \citet{Schaefer:2006pa}) ), we directly generate distance moduli like we do for SNe Ia for simplicity." +" The intrinsic scatter is set to be 0.65 magnitudes, which is approximately the average of the errors of the GRBs’ average distance moduli presented in ?,, and we ignore the measurement uncertainties, which are less than the intrinsic scatter."," The intrinsic scatter is set to be $0.65$ magnitudes, which is approximately the average of the errors of the GRBs' average distance moduli presented in \citet{Schaefer:2006pa}, and we ignore the measurement uncertainties, which are less than the intrinsic scatter." + We consider two kinds of distributions for GRBs in the redshift bin 1.8«z<7., We consider two kinds of distributions for GRBs in the redshift bin $1.8 < z < 7$. +" One is the uniform distribution, the other a very rough approximation to the distribution presented by Fig."," One is the uniform distribution, the other a very rough approximation to the distribution presented by Fig." +" 2 in ?,, i.e. P(z)«exp(—z/7)."," 2 in \citet{Bromm:2005ep}, i.e. $P(z) \propto \exp (-z/7)$." + We will see that our results are independent of the GRB distributions., We will see that our results are independent of the GRB distributions. + Figures | and 2 show our results for the constraints from GRBs distributed in the redshift bin 1.8 3\times +10^{-23}$, which is similar to the $1\sigma$ detection threshold of \citep[e.g., see recent discussions about AM~CVn binaries and {\it +LISA} ." + Iu summary. along with the original SDSS AM CVu found bv Roelofsetal.(2001.2005).. the £ additional new SDSS finds presented here provide a vield of 5 SDSS AM CVn candidates thus far: this is a substantial addition to the elite. AM. οδα subclass. compared to the dozen other cases previously known.," In summary, along with the original SDSS AM CVn found by \citet{roe04,roe05}, the 4 additional new SDSS finds presented here provide a yield of 5 SDSS AM CVn candidates thus far; this is a substantial addition to the elite AM CVn subclass, compared to the dozen other cases previously known." + Two of these five SDSS objects are also now strougly confirmed as nlizachort-period 1inanes (SDSS J1210-0159 bv Roclofs et al., Two of these five SDSS objects are also now strongly confirmed as ultrashort-period binaries (SDSS J1240-0159 by Roelofs et al. + 2005. aud J0926]|3621 iu this paper).," 2005, and J0926+3624 in this paper)." + SDSS J0926]|23621 reported here is the first coufideut example of au eclipsing AM. CVu., SDSS J0926+3624 reported here is the first confident example of an eclipsing AM CVn. + Our initial approximate considerations presented above for SDSS J0926|3621 presage that future detailed modeling of the 28.3 1umute eclipsing liebiteurve aud its long-terui teiiporal stability. double-peaked spectral line profiles. and follow-on radial velocity studies and multivaveleusthi observations. should provide an excellent ορπο] testbed of various models for AM. CV systems.," Our initial approximate considerations presented above for SDSS J0926+3624 presage that future detailed modeling of the 28.3 minute eclipsing lightcurve and its long-term temporal stability, double-peaked spectral line profiles, and follow-on radial velocity studies and multiwavelength observations, should provide an excellent empirical testbed of various models for AM CVn systems." +pixclization library.,pixelization library. + It cau inchide several subroutines aud operating programs., It can include several subroutines and operating programs. + The basic program of the second level. shown as a big rectangle. interacts with the first level subroutines.," The basic program ' of the second level, shown as a big rectangle, interacts with the first level subroutines." + These subroutines are shown bv small rectaneles aud call external Hibraries for the Foumier transtorm aud Legeudre polynomial calculations., These subroutines are shown by small rectangles and call external libraries for the Fourier transform and Legendre polynomial calculations. + The package reads and writes data both in ASCII table aud FITS formats., The package reads and writes data both in ASCII table and FITS formats. + More than 10 programs of the GLESP package operate in the GLESP zone., More than 10 programs of the GLESP package operate in the GLESP zone. + The prescut development of the package has also parallel calculation nuplemieutation., The present development of the package has also parallel calculation implementation. + Visualization procedures in GCL have been developed at IaO. Cambridge.," Visualization procedures in GL have been developed at IaO, Cambridge." + Three tests allow us to check the code., Three tests allow us to check the code. +" The first of them is from the analytical maps to calculate a;,,.", The first of them is from the analytical maps to calculate $a_{\ell m}$. + The code reproduces the theoretical Gon better than 10.*., The code reproduces the theoretical $a_{\ell m}$ better than $10^{-7}$. +" The secoud test is to reproduce an analytical map AG.0)2YuGeo) from a eiven ep, These tests check the caleulatious of the map and spherical cocficicuts incdependcutly."," The second test is to reproduce an analytical map $\Delta T(x,\phi)=Y_{\ell m}(x,\phi)$ from a given $a_{lm}$ These tests check the calculations of the map and spherical coefficients independently." + The third test is the reconstruction of αμ after the calculations of the map. ATGre.o). and back.," The third test is the reconstruction of $a_{\ell m}$ after the calculations of the map, $\Delta T(x,\phi)$, and back." + This test allows oue to check orthogonality., This test allows one to check orthogonality. + If the transformation is based ou really orthogonal functious it has to return after forward and backward calculation the same κ values., If the transformation is based on really orthogonal functions it has to return after forward and backward calculation the same $a_{\ell m}$ values. + Precision of the code cau be estimated by introduction of a set of τν=1 and recoustruction of them., Precision of the code can be estimated by introduction of a set of $a_{\ell m}=1$ and reconstruction of them. +" This test showed that using relation (13)) we can reconstruc the introduced 24;,, with the precision 10.* limited only by sinele precision of float point data recording aud with the precision 10.? for relation (11))."," This test showed that using relation \ref{legm}) ) we can reconstruct the introduced $a_{\ell m}$ with the precision $\sim +10^{-7}$ limited only by single precision of float point data recording and with the precision $\sim 10^{-5}$ for relation \ref{legl}) )." + Fie., Fig. + 7. demonstrates the accuracy of C5 caleulatious using HEALPix audGLESTI?., \ref{hp_compare} demonstrates the accuracy of $C_\ell$ calculations using HEALPix and. +. It should be noted that unlike the IIEALPix code. the GLESP method does not needed amy iteration for calculation of the αν coefficients and therefore is much faster.," It should be noted that unlike the HEALPix code, the GLESP method does not needed any iteration for calculation of the $a_{\ell m}$ coefficients and therefore is much faster." +" Our definition of the e;,, coefficients is exactly the same as in HEALDix as an estimator of the anisotropy power spectrui: Auv re-pixelization procedure will cause loss of information and thereby introduce uncertainties aud errors.", Our definition of the $a_{\ell m}$ coefficients is exactly the same as in HEALPix as an estimator of the anisotropy power spectrum: Any re-pixelization procedure will cause loss of information and thereby introduce uncertainties and errors. + The GLESP code has procedures for map re- based on two cliffercut methods iu the AT(0.0) domain: the first one consists in averaging," The GLESP code has procedures for map re-pixelization based on two different methods in the $\Delta T(\theta,\phi)$ –domain: the first one consists in averaging" +is no reason to expect that of those remaining will coutain the true value of the parameter.,is no reason to expect that of those remaining will contain the true value of the parameter. + This is because iustead: of suunuaime over all values of iy to eet a probability that exceeds.. we are siunmniug ouly over those values ercater than the detection threshold.," This is because instead of summing over all values of $\nS$ to get a probability that exceeds, we are summing only over those values greater than the detection threshold." + This results ina form of aud is ciscussed in detail im rets:eddington.., This results in a form of and is discussed in detail in \\ref{s:eddington}. + Tn a Bavesiau setting. probability is used f) quantity uncertainty in knowledge and iu this regard paralucters are typicalv viewed as random quantities.," In a Bayesian setting, probability is used to quantify uncertainty in knowledge and in this regard parameters are typically viewed as random quantities." + This distinction leads to amore intuitive interpretation of the credible interva1., This distinction leads to a more intuitive interpretation of the credible interval. + A credible interval at the £% level. for exa]e. is auv interval that contains the true value of the parameor EY of the time accoring fo its posterior ¢istribution. (," A credible interval at the $L$ level, for example, is any interval that contains the true value of the parameter $L$ of the time according to its posterior distribution. (" +See Park et al.,See Park et al. + 2008 for discussion on interval selection.), 2008 for discussion on interval selection.) +" Thus. from a Bayοσα perspective, it is proper to sav that there is au L'A chance that the source intensity is coutaimed iu t16 reported credible interval."," Thus, from a Bayesian perspective, it is proper to say that there is an $L$ chance that the source intensity is contained in the reported credible interval." + The corresponding credible intervals look similar to he confidence intervals iu Figure 1.. at least i1 Ligh countscenarios’.," The corresponding credible intervals look similar to the confidence intervals in Figure \ref{fig:CI}, at least in high count." + So far we have considered a very simple probenm with oulv oue unknown parameter. As.," So far we have considered a very simple problem with only one unknown parameter, $\lamS$." + The situation Is more complicated if trere are unknownparameters. such as Ap.," The situation is more complicated if there are unknown, such as $\lamB$." + In this case. frequency based intervals typically are constructed using asviiptotic arguments and/or by conditioning on ancillary statistics that vield a conditional sampling distribution tha does not depend on the nuisauce parameter.," In this case, frequency based intervals typically are constructed using asymptotic arguments and/or by conditioning on ancillary statistics that yield a conditional sampling distribution that does not depend on the nuisance parameter." + Ideutitving ancillary statistics can be a subtle task aud the resulting intervals may not be unique., Identifying ancillary statistics can be a subtle task and the resulting intervals may not be unique. + Bayesian intervals. can be coustructed using a simple and clear principle known as11argiualization.," Bayesian intervals, can be constructed using a simple and clear principle known as." + Forexample. if Ap is unknown. the mareinal posterior distribution of As is simply Credible intervals for Ag are computed just as before. but using the marginal posterior distribution.," Forexample, if $\lamB$ is unknown, the marginal posterior distribution of $\lamS$ is simply Credible intervals for $\lamS$ are computed just as before, but using the marginal posterior distribution." + Weeuphasize that neither confidence nor credible intervals directly quantity the detection seusitivitv of an experiment., We emphasize that neither confidence nor credible intervals directly quantify the detection sensitivity of an experiment. + To do this we cousicer the detecfiou problem in detail. which frou a statistical point of view is a test of the hypothesis tha there is no source emission in the given energy baud. ic. a test of," To do this we consider the detection problem in detail, which from a statistical point of view is a test of the hypothesis that there is no source emission in the given energy , i.e., a test of" +where Qíqgo) is the Heaviside step function.,"where $\Theta +{(y_2) }$ is the Heaviside step function." + Note that & is the caustic curvature at the origin which enters explicitly into the amplification formula.," Note that $\kappa$ is the caustic curvature at the origin \citep {Gaudi,alzh_03} + which enters explicitly into the amplification formula." + Formula (16)) vields an effective approximation for the point source magnification near the coordinate origin provided ju yo>0. and οντ is not too small (see the term containing &).," Formula \ref{point amplification}) ) yields an effective approximation for the point source magnification near the coordinate origin provided that $y_2> 0$, and ${y_2}/{y_1^2}$ is not too small (see the term containing $\kappa$ )." + For a fixed source position. this can be satisfied always » àn appropriate choice of the coordinate origin. so that the source will be situated almost on a normal to the tangent to 1ο CAUSLIC.," For a fixed source position, this can be satisfied always by an appropriate choice of the coordinate origin, so that the source will be situated almost on a normal to the tangent to the caustic." + If the source is on the caustic tangent or in the region between the caustic and the tangent. then formula (16)) does not represent a good approximation to the point source magnification.," If the source is on the caustic tangent or in the region between the caustic and the tangent, then formula \ref{point +amplification}) ) does not represent a good approximation to the point source magnification." + Nevertheless. in case of an extended source. we will show ju result. (16)) can be used to obtain approximations to the amplification of this source even as it intersects the caustic.," Nevertheless, in case of an extended source, we will show that result \ref{point amplification}) ) can be used to obtain approximations to the amplification of this source even as it intersects the caustic." + However. to do this. we need to redefine correctly the convolution of (16)) with a brightness distribution.," However, to do this, we need to redefine correctly the convolution of \ref{point amplification}) ) with a brightness distribution." + Let L(y) be a surface brightness distribution of an extended source., Let $I({\bmath y})$ be a surface brightness distribution of an extended source. + HE the source center is located at the point Y=(3.313) in the source plane. then the total microlensed Dux from the source is where the point source amplification A(y)=SoA; is the sum of amplifications of all the images.," If the source center is located at the point ${\bmath{Y}}=(Y_1,Y_2)$ in the source plane, then the total microlensed flux from the source is where the point source amplification $K( {\bmath y} ) = +\sum\limits_i {K_i } $ is the sum of amplifications of all the images." + The result of using the first integral [rom Iq. (18)), The result of using the first integral from Eq. \ref{flux_extended}) ) + obviously is equivalent to the result of the well-known ray-tracing method (when the pixel sizes tend to zero)., obviously is equivalent to the result of the well-known ray-tracing method \citep{schneider_92} (when the pixel sizes tend to zero). +" Near a caustic. one can approximate A(y)=Au|fis fy). where Au is an amplification ofall noncritical images that is supposed to be constant during LAL. and AG, is the amplification of the critical images."," Near a caustic, one can approximate $K( {\bmath y} )=K_0 + K_{cr} ( +{\bmath y})$ , where $K_0 $ is an amplification of all noncritical images that is supposed to be constant during HAE, and $K_{cr}$ is the amplification of the critical images." +" Formula (16)) contains the non-integrable term OCp)(yo)""7.", Formula \ref{point amplification}) ) contains the non-integrable term $\sim\Theta(y_{2}) ( {y_2 })^{ - 3 /2} $. + Therefore. the question arises of how formula (16)) can be used in situation when the extended. source intersects a caustic and some part of the source is in the zone between the tangent and the caustic.," Therefore, the question arises of how formula \ref{point amplification}) ) can be used in situation when the extended source intersects a caustic and some part of the source is in the zone between the tangent and the caustic." + In view of Section 3.2.. Ht is evident that the mentioned term is a result of the expansion of the root Viedyp/2| «in the approximate solution (11--14)).," In view of Section \ref{ss2.2}, it is evident that the mentioned term is a result of the expansion of the root $\sqrt +{y_2 + \kappa y_1^2 t^2/2+...}$ in the approximate solution \ref{new1}- \ref{new1Z}) )." + Any non-integrable terms in Avi; does not arise without using this expansion., Any non-integrable terms in $K_{cr}$ does not arise without using this expansion. +" Lt is easy to show that. in order to define A, correctly. one must replace the term OCyo)Cy)7/2 in (16)) bv the clistribution (generalized function) (jo)77 (Golfand&Shilov1964)."," It is easy to show that, in order to define $K_{cr}$ correctly, one must replace the term $\Theta(y_2) +( y_2 )^{ - 3 /2} $ in \ref{point amplification}) ) by the distribution (generalized function) $( y_2 )_+^{ - 3/2}$ \citep{Gel'fand_64}." +.. We recall that the distribution /57 of the variable Hy is defined by the expression [or any test function f(y)., We recall that the distribution $ y _+^{ - 3/2}$ of the variable $y$ is defined by the expression for any test function $f(y)$ . + After this redefinition. we have Ες formula can be used to correctly derive an approximate magnification of a sullicientlv smooth extended source including the case where the source crosses the caustic.," After this redefinition, we have This formula can be used to correctly derive an approximate magnification of a sufficiently smooth extended source including the case where the source crosses the caustic." + Now we use formula (19)) to derive the magnification of a Gaussian source with the brightness distribution where the parameter £L characterizes the source size., Now we use formula \ref{generalized Kcr}) ) to derive the magnification of a Gaussian source with the brightness distribution where the parameter $L$ characterizes the source size. +" The amplification of an extended source is defined as the ratio of the lensed Hux (18)) to the Hux of the unlensed source Ly=Ir1(g)dy,dyo which is equal to 1 in case of formula (200).The amplification of a Gaussian source Ive; is obtained by the substitution of (200) and (19)) into (18))"," The amplification of an extended source is defined as the ratio of the lensed flux \ref{flux_extended}) ) to the flux of the unlensed source $F_0 = \int\!\!\!\int I \left( {\rm {\bmath y}} +\right)dy_1 dy_2 $ which is equal to 1 in case of formula \ref{gaussian distribution}) ).The amplification of a Gaussian source $K_G$ is obtained by the substitution of \ref{gaussian distribution}) ) and \ref{generalized Kcr}) ) into \ref{flux_extended}) )." +" Further. we introduce thedimensionless coordinates s=Y,/L.hYS/L of the source centre and the functions"," Further, we introduce thedimensionless coordinates $s = Y_1/L,\, h += Y_2/ L$ of the source centre and the functions" +properties.,properties. +" In order to reproduce the observed properties the mergers are required to happen at high-redshift (z> 1), between progenitors of different mass ratio (at least 3:1) and with a significant fraction of the total mass in the form of gas (Le., >10 percent)."," In order to reproduce the observed properties the mergers are required to happen at high-redshift $z \geq 1$ ), between progenitors of different mass ratio (at least 3:1) and with a significant fraction of the total mass in the form of gas (i.e., $\geq 10$ percent)." +" However, this formation scenario is not able to recreate the the observed mass-metallicity gradient relation of these low-luminosity galaxies etal. 2009a;; Paper II)."," However, this formation scenario is not able to recreate the the observed mass-metallicity gradient relation of these low-luminosity galaxies \citealt{spolaor09a}; Paper II)." +" A further interpretation of our results is that low-luminosity galaxies were originally late-type galaxies, whose star formation has been truncated by removal of gas (ie., strangulation) and subsequently the disc has been dynamically heated by high velocity encounters (i.e., galaxy harassment) in the cluster environment."," A further interpretation of our results is that low-luminosity galaxies were originally late-type galaxies, whose star formation has been truncated by removal of gas (i.e., strangulation) and subsequently the disc has been dynamically heated by high velocity encounters (i.e., galaxy harassment) in the cluster environment." + Simulations have shown that late-type galaxies entering in a rich cluster can undergo a significant morphological transformation into spheroidals by encounters with brighter galaxiesand the cluster's tidal field (Mooreetal.1996;; Mastropietroetal. 2005))., Simulations have shown that late-type galaxies entering in a rich cluster can undergo a significant morphological transformation into spheroidals by encounters with brighter galaxiesand the cluster's tidal field \citealt{moore96}; \citealt{mastropietro05}) ). +" In this scenario, we expect the original disc to be dynamically heated by the interactions such that stellar orbits acquire a significant velocity component perpendicular to the disc."," In this scenario, we expect the original disc to be dynamically heated by the interactions such that stellar orbits acquire a significant velocity component perpendicular to the disc." + The imprint of their previous morphological nature is preserved in the form of an embedded stellar disc (e.g.; DeRijckeetal.2005: Chilingarianetal. 2008)).," The imprint of their previous morphological nature is preserved in the form of an embedded stellar disc (e.g., \citealt{rijcke03}; \citealt{chilingarian08}) )." +" For example, Beasleyetal.(2009) reported significant rotation at large radii (i.e., 4—7r.) of two luminous Virgo dwarf ellipticals, using globular cluster systems as tracers of galaxy dynamics."," For example, \cite{beasley09} reported significant rotation at large radii (i.e., $4 -7 r_{e}$ ) of two luminous Virgo dwarf ellipticals, using globular cluster systems as tracers of galaxy dynamics." + They show that the detection of such large amount of rotation in the outer galactic regions support the idea that luminous dwarf ellipticals were originally disc galaxies., They show that the detection of such large amount of rotation in the outer galactic regions support the idea that luminous dwarf ellipticals were originally disc galaxies. +" Numerical simulations (Mastropietroetal. 2005)) predict values of the anisotropy parameter (υ/σ)” similar to those of our galaxies, and also disky isophotes with similar B4 values."," Numerical simulations \citealt{mastropietro05}) ) predict values of the anisotropy parameter $(v/\sigma)^{*}$ similar to those of our galaxies, and also disky isophotes with similar $\overline{B_{4}}$ values." +" Moreover, the lack of counter rotation and the high incidence of coupled rotation between disc and bulge observed in our sample galaxies may favour a disc-heating scenario, whereby stars of the original disc contribute towards the bulge population while retaining some of the angular momentum."," Moreover, the lack of counter rotation and the high incidence of coupled rotation between disc and bulge observed in our sample galaxies may favour a disc-heating scenario, whereby stars of the original disc contribute towards the bulge population while retaining some of the angular momentum." +" However, this scenario is more difficult to reconcile with the stellar population properties, and in particular the mass-metallicity gradient relation, observed in our galaxies."," However, this scenario is more difficult to reconcile with the stellar population properties, and in particular the mass-metallicity gradient relation, observed in our galaxies." + We expect the star formation to be truncated in the disc due to the interactions., We expect the star formation to be truncated in the disc due to the interactions. +" Thus, metallicity gradients are required to form in the late-type galaxies and to be somehow preserved during the high velocity encounters."," Thus, metallicity gradients are required to form in the late-type galaxies and to be somehow preserved during the high velocity encounters." +" We conclude that although the kinematic and isophotal features in our galaxy sample can be interpreted in the context of a morphological transformation from late to early types, further detailed numerical simulations are needed to understand if such a scenario can also explain the stellar population trends observed."," We conclude that although the kinematic and isophotal features in our galaxy sample can be interpreted in the context of a morphological transformation from late to early types, further detailed numerical simulations are needed to understand if such a scenario can also explain the stellar population trends observed." + We have investigated the kinematic and photometric properties at large galactocentric radii for a sample of 14 low-luminosity early-type galaxies., We have investigated the kinematic and photometric properties at large galactocentric radii for a sample of 14 low-luminosity early-type galaxies. +" The radial extent considered in our analysis, ie. ~1— 3re, allows us to"," The radial extent considered in our analysis, i.e. $\sim 1 - 3 r_{e}$ , allows us to" +with the ereater rate of bGE.2) is dominant all the wav from redshift z* to :.,"with the greater rate of $b(E,z)$ is dominant all the way from redshift $z^\star$ to $z$." +" The transfer equation is solved (Montiuerle1977). to obtain the CCR energy spectruui from a CR burst at τς, where O;uQE.=@AL.lgjego) is) the normalized flux τνof τν)/ per comoving volune with Q;(E.:.:,)=INE. i). 3 and 3” are the velocities corresponding to enerev E aud £%. respectively."," The transfer equation is solved \citep{mon1977} to obtain the CCR energy spectrum from a CR burst at $z_s$, where $\Phi_{i,{\rm H}}(E,z,z_s)\equiv \Phi_i(E,z,z_s)/n_{\rm H}(z)$ is the normalized flux of $i$ per comoving volume with $\Phi_i(E,z,z_s)\equiv +\beta N_i(E,z)_{z_s}$ , $\beta$ and $\beta'$ are the velocities corresponding to energy $E$ and $E'_s$, respectively." + μη] is the present average nmuuber density of protons iu the universe., $n_{\rm H}^0$ is the present average number density of protons in the universe. + © is an effect resulting when the nuclear destruction is considered. aud given as After analysis with Eq. (13)).," $\xi$ is an effect resulting when the nuclear destruction is considered, and given as After analysis with Eq. \ref{eq13}) )," + oue can estimate l9:ΟΙ|epomELS]tldsfeel.1o aud find au expression for c; The production rate of ποτ clement / of energy. E. produced at redshift 2 is elven by where oj);4CE.EY) is a cross section of a process between a CR unclide / with energv per uncleon E aud a backeround species j to make a eiven liebt clement 7 with E. aud »j(:) aud D) are background nunber abuudance of a uuclide j aud KEENnuuber ratio ofj to proton. respectively.," one can estimate $|\partial +z^\star / \partial E'|_{E'=E'_s}=|b(E'_s,z_s)|^{-1} |dz/dt|_{z=z_s}$ and find an expression for $\Phi_{i, {\rm H}}$ The production rate of light element $l$ of energy $E$ , produced at redshift $z$ is given by where $\sigma_{ij \rightarrow l}(E,E')$ is a cross section of a process between a CR nuclide $i$ with energy per nucleon $E'$ and a background species $j$ to make a given light element $l$ with $E$, and $n_j(z)$ and $K_{jp}^{\rm IGM}(z)$ are background number abundance of a nuclide $j$ and number ratio of $j$ to proton, respectively." + When the destruction of the liebt clement / after production is neelected. the total production rate is calculated as wheregit7j|GE!) is the total cross section of a reaction i|po>X. with anv OX.," When the destruction of the light element $l$ after production is neglected, the total production rate is calculated as where $\sigma_{ij \rightarrow l}^{\rm tot}(E')$ is the total cross section of a reaction $i+j \rightarrow l+X$, with any $X$." +" I adopt cross sections from ReadViola(1981).. aud particularly for the a|α reaction. exponcutial-plus-constaut cross section for | —9Li aud exponential one for { =""Li trom Merceretal. (2001)."," I adopt cross sections from \citet{rea1984}, and particularly for the $\alpha+\alpha$ reaction, exponential-plus-constant cross section for $l=^6$ Li and exponential one for $l=^7$ Li from \citet{mer2001}." +. The resulting light. clement abundauce is obtained as the CR production added to the BBN vield., The resulting light element abundance is obtained as the CR production added to the BBN yield. + The vield by CR uncleosvuthesis is the integration of those produced at z/ πο CRs generated at z4 over +! ando. thus I also caleulate the. LiDoD production i the universe by the secoudarv process. L6. [pn]es [CO]imar >[LiBeBlisar.," The yield by CR nucleosynthesis is the integration of those produced at $z'$ from CRs generated at $z_s$ over $z'$ and $z_s$, thus I also calculate the LiBeB production in the universe by the secondary process, i.e., $p\alpha$ $_{\rm CR}$ $_{\rm +ISM}\rightarrow$ $_{\rm ISM}$." + Since the C. and— ο abundances of the ISALD in structures are about two orders of magnitude higher than those of the ICAL (seeFie.1linDaignueetal.2006).. the secondary LiBeB production in the ICAL is uot important.," Since the C and O abundances of the ISM in structures are about two orders of magnitude higher than those of the IGM \citep[see Fig. 11 in][]{dai2006}, the secondary LiBeB production in the IGM is not important." + I expect that the LiBeB abundances in the ISAL are enhanced b a contribution of the secondary process., I expect that the LiBeB abundances in the ISM are enhanced by a contribution of the secondary process. + Iu fact. the reactions of [paler Ομ»[LiDeD]i make light elements in the ISALD anc the mass acerction to the structures from the IGAL dilutes the ISM. abuudanuces in the framework of this model involviug a hierarchical structure formation.," In fact, the reactions of $p\alpha$ $_{\rm CR}$ $_{\rm +ISM}\rightarrow$ $_{\rm ISM}$ make light elements in the ISM and the mass accretion to the structures from the IGM dilutes the ISM abundances in the framework of this model involving a hierarchical structure formation." + Note that from the assumption that the confinement of CRs by a iiagnuetic field is ineffective. the CRs do not stav in the structures.," Note that from the assumption that the confinement of CRs by a magnetic field is ineffective, the CRs do not stay in the structures." + The πο clement abundances produced by the secondary reactions are then eiven with a parameter: the fraction of barvous at redshift + which are in structures where fpaQM.2) is the distribution function of halos taken from the Sheth&Tormen(1999) inodification the Press-Schechter function (Press&Schechter1971L) converted iuto the mass function (Jemlksinsetal.2001) bv a code provided Jn A. Jenkins (2007. priva5 conmuimnication).," The light element abundances produced by the secondary reactions are then given with a parameter: the fraction of baryons at redshift $z$ which are in structures where $f_{\rm PS}(M,z)$ is the distribution function of halos taken from the \citet{she1999} modification to the Press-Schechter function \citep{pre1974} converted into the mass function \citep{jen2001} by a code provided by A. Jenkins (2007, private communication)." + Tassie that the primordial power spectral slope is (—1. the rus amplitucle for mass deusity fluctuations iu a sphere of radius 8 5.1 Mpc is oy=0.9. and the Boud&Efstathiou(1981). fit to the trauster function for cold dark matter is used in geuerating a nass function.," I assume that the primordial power spectral slope is $n$ =1, the rms amplitude for mass density fluctuations in a sphere of radius 8 $h^{-1}$ Mpc is $\sigma_8=0.9$, and the \citet{bon1984} fit to the transfer function for cold dark matter is used in generating a mass function." + ppar is the comoving dark matter deusitv of the Undverse., $\rho_{\rm DM}$ is the comoving dark matter density of the universe. + The light elements made by the secondary process are contained in the structuresthat erow eraduallv., The light elements made by the secondary process are contained in the structuresthat grow gradually. + The abundance of a light clement inthe structures is then eiven bv I calculate the light clement production iu the uuiform universe by CCRs (ic. neglecting an mhbomosgeneity of CCTx)., The abundance of a light element inthe structures is then given by I calculate the light element production in the uniform universe by CCRs (i.e. neglecting an inhomogeneity of CCRs). + T consider only processes between the accelerated CRs with the abundance patterns of the structures iu the Daigueetal.(2006) model aud the backerouud, I consider only processes between the accelerated CRs with the abundance patterns of the structures in the \citet{dai2006} model and the background +At higher X-ray energies the boundary of the emission zone perpendicular to the jet direction is rather sharp. the intensity increases typically by a factor of ~ 2.5 and shows a flat top profile.,"At higher X-ray energies the boundary of the emission zone perpendicular to the jet direction is rather sharp, the intensity increases typically by a factor of $\sim$ 2.5 and shows a flat top profile." + At low energies the profile smears out and widens. as if the radiating electrons diffuse out of the acceleration zone.," At low energies the profile smears out and widens, as if the radiating electrons diffuse out of the acceleration zone." + This can be seen in reffig:;jetl-s-h.., This can be seen in \\ref{fig:jet1-s-h}. + In the hard 2-12 kkeV energy band (bottom figure) the emission region is sharply bounded as explained above and clearly pointing towards SS4433., In the hard $-$ keV energy band (bottom figure) the emission region is sharply bounded as explained above and clearly pointing towards 433. + The soft O0.2-]kkeV image (top) shows only an extended diffuse emission region. centered on the lenticular structure. without any connection to on the right. outside the images.," The soft $-$ keV image (top) shows only an extended diffuse emission region, centered on the lenticular structure, without any connection to on the right, outside the images." + The transfer of kinetic energy of the relativistically flowing matter into radiation is obviously a complex process and it is not quite clear whether it is controlled “intrinsically” (i.e. just by the flow itself). or whether environmental conditions play a role as well.," The transfer of kinetic energy of the relativistically flowing matter into radiation is obviously a complex process and it is not quite clear whether it is controlled `intrinsically' (i.e. just by the flow itself), or whether environmental conditions play a role as well." +" Note that the outer shape of the surrounding radio remnant W50 shows a very pronounced northern bay at a similar right ascension as the X-ray emission maximum (~ 19"" 14”). which does not have a counterpart on the southern side."," Note that the outer shape of the surrounding radio remnant W50 shows a very pronounced northern bay at a similar right ascension as the X-ray emission maximum $\sim$ $^h$ $^m$ ), which does not have a counterpart on the southern side." + At the outermost eastern boundary of W50 the X-ray and the radio structures show an astonishingly similar appearance. indicating the interaction of a “terminal shock’ of the jet and the circumstellar medium.," At the outermost eastern boundary of W50 the X-ray and the radio structures show an astonishingly similar appearance, indicating the interaction of a `terminal shock' of the jet and the circumstellar medium." + The interior of the radio remnant is filled with X-ray emitting material. and the spatial coincidence of X-ray and radio emission suggests that the physical conditions of the terminal shock region are very similar to those found at the outer shocks of ‘ordinary’ supernova remnants.," The interior of the radio remnant is filled with X-ray emitting material, and the spatial coincidence of X-ray and radio emission suggests that the physical conditions of the terminal shock region are very similar to those found at the outer shocks of 'ordinary' supernova remnants." + Interestingly. the prominent X-ray emission seems to form a ring-like structure confined by the outer boundary of the radio remnant.," Interestingly, the prominent X-ray emission seems to form a ring-like structure confined by the outer boundary of the radio remnant." + This indicates that the final jet flow might still have some kind of hollow-cone morphology., This indicates that the final jet flow might still have some kind of hollow-cone morphology. + XMM-Newton has for the first time the sensitivity and the wide band pass to study in more detail the spectral behavior of the precessing jet interacting with its ambient medium., XMM-Newton has for the first time the sensitivity and the wide band pass to study in more detail the spectral behavior of the precessing jet interacting with its ambient medium. + However. the photon statistics still do not allow accurate fits of individual small spatial structures visible in the data.," However, the photon statistics still do not allow accurate fits of individual small spatial structures visible in the data." + Fortunately. there are no indications for drastic spectral changes occuring on small spatial scales (e.g.. Figs 3 and 4).," Fortunately, there are no indications for drastic spectral changes occuring on small spatial scales (e.g., Figs 3 and 4)." + A simple power law model provides an excellent representation of the spectrum of the main bright lenticular area of reffig:rgbl.., A simple power law model provides an excellent representation of the spectrum of the main bright lenticular area of \\ref{fig:rgb1}. + The obtained slope of [=2.17+0.02 1s in the range commonly found for non-thermal supernova remnants like in the Crab., The obtained slope of $\Gamma = 2.17\pm0.02$ is in the range commonly found for non-thermal supernova remnants like in the Crab. + A similar power law is seen in the outer. terminal shock region although there the low statistical significance of the high energy data does not allow a very precise determination of the slope.," A similar power law is seen in the outer, terminal shock region although there the low statistical significance of the high energy data does not allow a very precise determination of the slope." + It must further be noticed that the bright spot can be equally well fitted with a bremsstrahlung model with kT = 3.9820. kkeV. a temperature in agreement with previous findings from ASCA (Yamauchi et al.," It must further be noticed that the bright spot can be equally well fitted with a bremsstrahlung model with kT = $\pm$ keV, a temperature in agreement with previous findings from ASCA (Yamauchi et al." + 1994). ROSAT ll) and Einstein (Watson 1983).," 1994), ROSAT I) and Einstein (Watson 1983)." +" However. we tend to believe that this model is merely a ""statistically acceptable? numerical representation of the spectral shape and not a correct description of the physical emission conditions. in particular. as the data show no indications of kkeV iron line emission. which would be expected at these temperatures."," However, we tend to believe that this model is merely a 'statistically acceptable' numerical representation of the spectral shape and not a correct description of the physical emission conditions, in particular, as the data show no indications of keV iron line emission, which would be expected at these temperatures." + In both cases. power law or simple bremsstrahlung. we find in the soft band spectral residuals resembling well known emission lines.," In both cases, power law or simple bremsstrahlung, we find in the soft band spectral residuals resembling well known emission lines." + While in the lenticular spot these features are relatively weak they are clearly dominant in the terminal shock region., While in the lenticular spot these features are relatively weak they are clearly dominant in the terminal shock region. + Adding to the power law a model resulted 1αυ] temperatures of ~0.2—0.3 kkeV for the thermal models. but the residuals still showed some line features which were not accounted for in this thermal tonization equilibrium model.," Adding to the power law a model resulted in temperatures of $\sim 0.2-0.3$ keV for the thermal models, but the residuals still showed some line features which were not accounted for in this thermal ionization equilibrium model." + Using the non-equilibrium model instead of improved the quality of the fits without significant change to the basic physical parameters., Using the non-equilibrium model instead of improved the quality of the fits without significant change to the basic physical parameters. + For the lenticular region the temperature remained at KT = 0.22+0.01 kkeV. for the terminal shock we obtained a kT = 0.35-0.01 kkeV: in both cases," For the lenticular region the temperature remained at kT = $\pm$ keV, for the terminal shock we obtained a kT = $\pm$ keV; in both cases" +measured with the All-Sky. Monitor. (ASAI) (Levinectal.1996) on the RNTE.,measured with the All-Sky Monitor (ASM) \cite{Levine96} on the RXTE. + NPE 01112το]τς location was in the field-of-view of the RNTE PCA (Jahodaetal.1996) and Ligh LEnerey Timing Experiment (LIENTE) (Rothschileletal.10905] for 42 observations [rom 1998 October 20 - 1999 May H1., XTE J0111.2–7317's location was in the field-of-view of the RXTE PCA \cite{Jahoda96} and High Energy Timing Experiment (HEXTE) \cite{Rothschild98} for 42 observations from 1998 October 20 - 1999 May 11. + Further. the Burst and Transient Source Experiment (BATSE) (Pishmanetal.1989). on the Compton Gamma Rav Observatory had. also observed NTEJOIII.27317's location from 1991-2000.," Further, the Burst and Transient Source Experiment (BATSE) \cite{Fishman89} on the Compton Gamma Ray Observatory had also observed XTE J0111.2–7317's location from 1991-2000." + No accditional outbursts have been detected with either BATSE or RATE., No additional outbursts have been detected with either BATSE or RXTE. + To determine pulse frequencies for NTE JOLIL.27317. we performed a grid search over a range of candidate frequencies using I-second. resolution 20-50 keV. data from the BATSE Large Area Detectors (LADs).," To determine pulse frequencies for XTE J0111.2–7317, we performed a grid search over a range of candidate frequencies using 1-second resolution 20-50 keV data from the BATSE Large Area Detectors (LADs)." + This technique is described in detail elsewhere. (Fingereal.1999:Wilson-LlodgeWilsonetal.2002: 2003).," This technique is described in detail elsewhere \cite{Finger99,Wilsonhodge99,Wilson02,Wilson03}." +. First we combined the count rates over the 4 LADs viewing NTIJ0111.2.1317. using weights optimized for àn exponential energy spectrum. {ας=Aexp(Lyk) witi temperature AL=12 keV. and then grouped them into 300-5J segments.," First we combined the count rates over the 4 LADs viewing XTE J0111.2–7317, using weights optimized for an exponential energy spectrum $f(E) = A \exp(-E/kT)$ with temperature $kT = 12$ keV, and then grouped them into 300-s segments." + In cach 300-s segment. we fitted a model consisting of a sixth-order Fourier expansion in pulse phase model (representing the 20-50 keV pulse profile) and a spline function with quaclratics in time (representing the background).," In each 300-s segment, we fitted a model consisting of a sixth-order Fourier expansion in pulse phase model (representing the 20-50 keV pulse profile) and a spline function with quadratics in time (representing the background)." + Using 300-5 segments allowed: use of a quadratic background. model. allowed: us o ellectively fit bright Earth occultation steps in the data. improved computational efficiency. and reduced smearing of he pulse profiles if the initial phase mocel was incorrect.," Using 300-s segments allowed use of a quadratic background model, allowed us to effectively fit bright Earth occultation steps in the data, improved computational efficiency, and reduced smearing of the pulse profiles if the initial phase model was incorrect." + The final step was to perform a grid search in. frequency using the set of typically several hundred. estimated 20-50 keV. pulse profiles in cach αν interval of data., The final step was to perform a grid search in frequency using the set of typically several hundred estimated 20-50 keV pulse profiles in each 4-day interval of data. + For each erid point. the segment pulse profiles were shifted according he the corresponding frequency. olfset.," For each grid point, the segment pulse profiles were shifted according the the corresponding frequency offset." + These profiles were hen combined into a mean profile., These profiles were then combined into a mean profile. +" Phe best. [frequency was selected using a modilied. ZZ. statistic. called Y, after (Finger""οἱal.1999).. given. by Y,.=A""lan°σαa2 where a, is the mean Fourier coellicient for harmonic f aud στah is the sample variance [from the segment profiles."," The best frequency was selected using a modified $Z_n^2$ statistic, called $Y_n$ after \cite{Finger99}, given by $Y_n = +\sum_{h=1}^n |\bar \alpha_h|^2/\sigma_{\bar \alpha_h}^2$ where $\bar +\alpha_h$ is the mean Fourier coefficient for harmonic $h$ and $\sigma_{\bar \alpha_h}^2$ is the sample variance from the segment profiles." + Using sample variances from the segment profiles rather than assuming Poisson statistics properly accounted for aperiodic noise in NTE J0111.2-7317 and in any other sources in the large BATSE field of view., Using sample variances from the segment profiles rather than assuming Poisson statistics properly accounted for aperiodic noise in XTE J0111.2-7317 and in any other sources in the large BATSE field of view. + A similar technique was used to generate pulse frequency measurements for RATE PCA observations [rom 2-60 keV Standard 1 (125 ms. no energy resolution) data and from. RATE LIENTE data.," A similar technique was used to generate pulse frequency measurements for RXTE PCA observations from 2-60 keV Standard 1 (125 ms, no energy resolution) data and from RXTE HEXTE data." + The root-mean-squarecd (rms). pulsed. fluxes. were estimated at the M⋅⋅ ↴lou∣⋡∢⋅≱∖∐∣⇂↓⋅⋖⋅⊏↥⋯⊾⊔≼⇍∙∖⇁⋜↧⊳∖∫↰↓⋩↼∖↓⋝∖∶∖∪⋅⋅↱≻“yh," The root-mean-squared (rms) pulsed fluxes were estimated at the best fit frequency as $F_{\rm RMS} = (0.5 \sum_{h=1}^m |\bar +\alpha_h|^2)^{1/2}$." + Pieure 4 shows the harveentered pulse frequeney history. measured with BATSLE (squares) ancl with the RATE PCA (open circles). the 20-50 keV. pulsed. (lux measured with BATSE. and the 2-60. keV. pulsed. Lux measured. with the RAPE PCA.," Figure \ref{fig:xr_hist} shows the barycentered pulse frequency history measured with BATSE (squares) and with the RXTE PCA (open circles), the 20-50 keV pulsed flux measured with BATSE, and the 2-60 keV pulsed flux measured with the RXTE PCA." + Pulsations were detected with BATSE from 1905 October 30 - 1999 January 9., Pulsations were detected with BATSE from 1998 October 30 - 1999 January 9. + Phe RXTE PCA continued to detect pulsations until observations ceased on 1999 February 9., The RXTE PCA continued to detect pulsations until observations ceased on 1999 February 19. + Pulsations were not detected in two subsequent. PCA observations on 1999 March 26 and. 1999 Alay 11., Pulsations were not detected in two subsequent PCA observations on 1999 March 26 and 1999 May 11. + Pulse frequeney derivatives were computed w dillerencing adjacent BATSE pulse [reeuenevy measurements., Pulse frequency derivatives were computed by differencing adjacent BATSE pulse frequency measurements. + Figure 5 shows the correlation between pequency derivative and 20-50. Κον pulsed. Εαν.," Figure \ref{fig:xr_fdotvsflux} + shows the correlation between frequency derivative and 20-50 keV pulsed flux." + This correlation is well fitted with a power law with index 90.1. which is consistent with the index of 6/7 expected rom simple accretion theory. if acerction from a disk. is assumed.," This correlation is well fitted with a power law with index $0.9 \pm 0.1$, which is consistent with the index of $6/7$ expected from simple accretion theory, if accretion from a disk is assumed." + If a reliable bolometric correction can be derived or a source. this correlation can be usec to constrain he distance or the magnetico field strengthe if the other is known.," If a reliable bolometric correction can be derived for a source, this correlation can be used to constrain the distance or the magnetic field strength if the other is known." + Unfortunately in this case. large changes in the ruse fraction over the course of the outburst. prevented. us rom ceriving a reliable bolometric [ux (see Figure 6).," Unfortunately in this case, large changes in the pulse fraction over the course of the outburst prevented us from deriving a reliable bolometric flux (see Figure 6)." + Analvsis of the energy spectrum of NTE J0111.2.1317 was, Analysis of the energy spectrum of XTE J0111.2–7317 was +"consequence ofa pathological choice of calibrators: he long period Cepheids (P> 9"") in the sample all lie ou the cool side of the instability. strip aud have overly simall uuinosities. aud the short iod. Cepheids all lic towards the ceuter or hot side of the instability strip and have overly larec tmuinositics.","consequence of a pathological choice of calibrators: the long period Cepheids $P \ge 9^{\rm d}$ ) in the sample all lie on the cool side of the instability strip and have overly small luminosities, and the short period Cepheids all lie towards the center or hot side of the instability strip and have overly large luminosities." + The sclection of Cepheid calibrators or the PL relation clearly affects the results obtained., The selection of Cepheid calibrators for the PL relation clearly affects the results obtained. + An important factor affecting such an analysis is the reddening of individual Cepheids., An important factor affecting such an analysis is the reddening of individual Cepheids. + (ναι that Wesenheit formulations are generally designed to correct for the effects of interstellar reddening. they would appear to be the iustrunenut of choice for studies of C'ephlek istances.," Given that Wesenheit formulations are generally designed to correct for the effects of interstellar reddening, they would appear to be the instrument of choice for studies of Cepheid distances." + There is an alternative available using huuinosities directly. as demonstrated here. but it is difficult to apply such a rolatiouship in distaut galaxies where ouly visible and ueariufrared plotometry is typically available for member Cepheids.," There is an alternative available using luminosities directly, as demonstrated here, but it is difficult to apply such a relationship in distant galaxies where only visible and near-infrared photometry is typically available for member Cepheids." +" A Wesenheit formulation can be derived from| ιο AA data of Table adjusted dy —0.01 to account for the need of a possible zero-point shift in the calibrated zero-age main sequence, and the AST parallax data in Table 1.. as shown in Fig. 9.."," A Wesenheit formulation can be derived from the $M_V$ data of Table \ref{tab3}, adjusted by –0.04 to account for the need of a possible zero-point shift in the calibrated zero-age main sequence, and the parallax data in Table \ref{tab4}, as shown in Fig. \ref{fig9}." + For this purpose a color term of 3.3 was adopted (seeMadore&Freedinan 1991)., For this purpose a color term of 3.3 was adopted \citep[see][]{mf91}. +. A nou-paraimctric fit to the data produces a relationship described by: As a test of the validitv of the resulting relationship for extragalactic Cepleids. au analysis was made using the data frou Macrietal.(2006) for Cepheids in the inner regions of NGC 1258. where ietallicitics appear to be close to the solar values typical of nearby Galactic calibrators.," A non-parametric fit to the data produces a relationship described by: As a test of the validity of the resulting relationship for extragalactic Cepheids, an analysis was made using the data from \citet{mc06} for Cepheids in the inner regions of NGC 4258, where metallicities appear to be close to the solar values typical of nearby Galactic calibrators." + The availableDV data for the ealaxws Cepheids (Maciietal.2006) were analyzed usine the Wesenheit formulation. as shown in Fie. 10..," The available data for the galaxy's Cepheids \citep{mc06} were analyzed using the Wesenheit formulation, as shown in Fig. \ref{fig10}." + Although the scatter is lareer than for Calactic calibrators. the agreement is otherwise excellent. particularly the trend with period. which depends heavily upon the Weseuheit formulation and the validity of the Calactic calibration over all pulsation periods.," Although the scatter is larger than for Galactic calibrators, the agreement is otherwise excellent, particularly the trend with period, which depends heavily upon the Wesenheit formulation and the validity of the Galactic calibration over all pulsation periods." + Restricting thle sample to inner region Cepheids with P>129. as sugeested by Macrictal.(2006) απ a meas of clinunating bias arisiue from contaminated Cepheids. aud climinating a few outlicrs viclds a distance modulus for NGC [258 of 29.31£0.01. corresponding to a distance ο [4.26AOS Ape.," Restricting the sample to inner region Cepheids with $P \ge 12^{\rm d}$, as suggested by \citet{mc06} as a means of eliminating bias arising from contaminated Cepheids, and eliminating a few outliers yields a distance modulus for NGC 4258 of $29.31\pm0.04$, corresponding to a distance of $7.27\pm0.13$ Mpc." + Tha value agrees closcly with the geometrical distance estimate of 7.24l.1 Alpe obtained bv Ienustemetal.(1999) using ΟΠ masers., That value agrees closely with the geometrical distance estimate of $7.2\pm1.4$ Mpc obtained by \citet{he99} using OH masers. + Cepheids in the outer regions of NGC 12558 vield a distance modulus of 29.51+0.05. the (P42 dqiffereuce relative to the value obtained for Cepheids wing in the inner regious of NGC 1258 prestunablybeing tied to the lower metallicity of outer region Cepheids.," Cepheids in the outer regions of NGC 4258 yield a distance modulus of $29.51\pm0.05$, the $0^{\rm m}.2$ difference relative to the value obtained for Cepheids lying in the inner regions of NGC 4258 presumablybeing tied to the lower metallicity of outer region Cepheids." +Britzen et al. (,Britzen et al. ( +"1999, 2001) and Britzen (2002).","1999, 2001) and Britzen (2002)." + A statistical analysis of 5 GHz VLBI polarimetry data from 177 sources in the CJF survey has been presented by Pollack et al., A statistical analysis of 5 GHz VLBI polarimetry data from 177 sources in the CJF survey has been presented by Pollack et al. + The results presented here. serve as the basis for the analysis of the kinematics of the sources which will be discussed in. Britzen et al.," The results presented here, serve as the basis for the analysis of the kinematics of the sources which will be discussed in Britzen et al." + 2007a (hereafter Paper ID., 2007a (hereafter Paper II). + In Britzen et al., In Britzen et al. + 2007b (hereafter Paper U1) we present a correlation analysis between soft X-ray and VLBI properties of the CJF The CJF was designed as a state-of-the-art survey providing multiple epochs of VLBI observations of a large. complete sample.," 2007b (hereafter Paper III) we present a correlation analysis between soft X-ray and VLBI properties of the CJF The CJF was designed as a state-of-the-art survey providing multiple epochs of VLBI observations of a large, complete sample." + To date. it is indeed the largest multi-epoch survey of motions 1 terms of the number of sources anc jet components tracked.," To date, it is indeed the largest multi-epoch survey of motions in terms of the number of sources and jet components tracked." + In addition. it provides angular resolution and dynamic range appropriate to identify anc trace individual jet components reliably across the epochs.," In addition, it provides angular resolution and dynamic range appropriate to identify and trace individual jet components reliably across the epochs." + Due to its completeness. statistical statements can be made concerning the distributions of velocities. bending. patter motions. and changes in the brightness of jet components ar= their dependence on the core separation.," Due to its completeness, statistical statements can be made concerning the distributions of velocities, bending, pattern motions, and changes in the brightness of jet components and their dependence on the core separation." + Given its size. completeness. and the range of source properties spanned. this database should be of great utility for statistical studies.," Given its size, completeness, and the range of source properties spanned, this database should be of great utility for statistical studies." + Care was taken to ensure homogeneity in the observing strategy. data reduction. and data analysis.," Care was taken to ensure homogeneity in the observing strategy, data reduction, and data analysis." + We hope to have produced a body of data that can be used to develop and test physical theories of active nuclei in ways that have not previously been possible., We hope to have produced a body of data that can be used to develop and test physical theories of active nuclei in ways that have not previously been possible. + Continued. VLBI observations of the CJF sources have been performed since 1990 (see Table 1)), Continued VLBI observations of the CJF sources have been performed since 1990 (see Table \ref{obs.tab}) ). + Subsamples were observed in several global VLBI observations and in VLBA snapshot runs at 6 em wavelength between March 1990 and December 2000., Subsamples were observed in several global VLBI observations and in VLBA snapshot runs at 6 cm wavelength between March 1990 and December 2000. + The VLBA snapshot runs of CJF sources started in 1998., The VLBA snapshot runs of CJF sources started in 1998. + The observational strategy was to observe the sources 8 times in 5.5 minute snapshots in each observing session and to record the data over 32 MHz total bandwidths broken up into four baseband channels. with | bit sampling.," The observational strategy was to observe the sources 8 times in 5.5 minute snapshots in each observing session and to record the data over 32 MHz total bandwidths broken up into four baseband channels, with 1 bit sampling." + The data were correlated in We aimed at a minimum of three epochs for every source since we found from experience that the unambiguous determination of the jet component position and motion requires at least three observing epochs spread over roughly four years (minimum time span is less than one year)., The data were correlated in We aimed at a minimum of three epochs for every source since we found from experience that the unambiguous determination of the jet component position and motion requires at least three observing epochs spread over roughly four years (minimum time span is less than one year). + These observations are now complete: the last epoch for a subsample of 34 sources was obtained in December 2000., These observations are now complete; the last epoch for a subsample of 34 sources was obtained in December 2000. + In Table 1 we list the correlator codes. dates. bandwidth. polarization information. antenna arrays. correlator. number and length of scans. and a reference for further information.," In Table \ref{obs.tab} we list the correlator codes, dates, bandwidth, polarization information, antenna arrays, correlator, number and length of scans, and a reference for further information." + To create a homogeneous. statistically valid database. we started a systematic (re-)analysis of all epochs for all sources obtained in the 1990s.," To create a homogeneous, statistically valid database, we started a systematic (re-)analysis of all epochs for all sources obtained in the 1990s." +" All sources and epochs of all ""old (data before 1998) and ""newly"" obtained data sets have beer analyzed in the same standardized way.", All sources and epochs of all “old” (data before 1998) and “newly” obtained data sets have been analyzed in the same standardized way. + Despite using a global array for the older epochs anc the VLBA for the new epochs. we aimed at obtaining similar observing conditions. calibratior methods (to ensure a reliable calibration of the sources. we included in each observing run at least one calibrator source. 3C 279. which was observed at similar (1.v) ranges). anc reduction. techniques.," Despite using a global array for the older epochs and the VLBA for the new epochs, we aimed at obtaining similar observing conditions, calibration methods (to ensure a reliable calibration of the sources, we included in each observing run at least one calibrator source, 3C 279, which was observed at similar $(u, v)$ ranges), and reduction techniques." + Calibration and. fringe-fitting were done using standard procedures in the Astronomical Image Processing System (AIPS. Greisen 1990).," Calibration and fringe-fitting were done using standard procedures in the Astronomical Image Processing System (AIPS, Greisen 1990)." + With automated mapping within (v.2.4b. Shepherd 1997). making use of the script¢ufomap. we obtained Clean maps of similar quality for all data sets.," With automated mapping within (v.2.4b, Shepherd 1997), making use of the script, we obtained Clean maps of similar quality for all data sets." + Extended sources were reprocessed using (within. difmap). to move the observation phase-center.," Extended sources were reprocessed using (within ), to move the observation phase-center." + In addition. we re-c/eaned with a different pixel size to map the extended emission reliably as well.," In addition, we ed with a different pixel size to map the extended emission reliably as well." + A critical analysis of the VLBI jet components and a comparison of jet properties requires a quantitative determination of the jet components’ features., A critical analysis of the VLBI jet components and a comparison of jet properties requires a quantitative determination of the jet components' features. + We therefore fitted Gaussian model components directly to the observed visibilities (real and imaginary parts) using. the Levenberg-Marquardt non-linear. least squares minimization technique (program within. difinap) to fit the brightness. sizes and positions of the individual Jet We modelfitted all sources at all epochs independently. starting from a point source. and using circular Gaussian components. with parameters flux density (5). position in Cartesian coordinates (vv). and. semi-diameter (M).," We therefore fitted Gaussian model components directly to the observed visibilities (real and imaginary parts) using the Levenberg-Marquardt non-linear least squares minimization technique (program within ) to fit the brightness, sizes and positions of the individual jet We modelfitted all sources at all epochs independently, starting from a point source, and using circular Gaussian components, with parameters flux density $S$ ), position in Cartesian coordinates $x,y$ ), and semi-diameter $M$ )." + The positions were later also converted to polar coordinates (7.4). as explained below.," The positions were later also converted to polar coordinates $r,\theta$ ), as explained below." + We used circular components since it turned out that the estimations of the axial ratio parameters were often enough ill-conditioned to make it unjustifiable to include them as free parameters in the model (1.e.. they are very highly correlated with other parameters).," We used circular components since it turned out that the estimations of the axial ratio parameters were often enough ill-conditioned to make it unjustifiable to include them as free parameters in the model (i.e., they are very highly correlated with other parameters)." + The sizes of the circular jet components were allowed to vary between epochs., The sizes of the circular jet components were allowed to vary between epochs. + We stopped adding jet components within the model-fitting process whenever a solution had been obtained. such that adding an additional component would not improve the quality of the fit. 1.8. reduce the value for chi-square. We calculated (statistical) uncertainties for the fitted Gaussian parameters for each source at each epoch via a slight modification todifmap.," We stopped adding jet components within the model-fitting process whenever a solution had been obtained, such that adding an additional component would not improve the quality of the fit, i.e. reduce the value for chi-square, We calculated (statistical) uncertainties for the fitted Gaussian parameters for each source at each epoch via a slight modification to." + We derived the covariance matrix. C. from the Hessian matrix (which was already computed during the model-fitting procedure) by using the pre-existing//-covar function. which simply performed Gauss-Jordan elimination to invert the matrix.," We derived the covariance matrix, $\mathcal{C}$, from the Hessian matrix (which was already computed during the model-fitting procedure) by using the pre-existing function, which simply performed Gauss-Jordan elimination to invert the matrix." +" The uncertainty for the 7"" parameter", The uncertainty for the $i^{\rm th}$ parameter +"for beam attenuation and pointing offsets, while it is unclear whether Garcia-Barreto et al. (","for beam attenuation and pointing offsets, while it is unclear whether Garcia-Barreto et al. (" +1994) used similar corrections.,1994) used similar corrections. +" An additional difference of 0.03 in HI—def is due to the fact that, to transform SHI fluxes into gas masses, Garcia-Barreto et al (1994) used in eq. ("," An additional difference of 0.03 in $HI-def$ is due to the fact that, to transform $SHI$ fluxes into gas masses, Garcia-Barreto et al (1994) used in eq. (" +1) a constant value of 2.22 x 10? instead of 2.36 x 10? as in this work.,1) a constant value of 2.22 $\times$ $^5$ instead of 2.36 $\times$ $^5$ as in this work. +" The relationship between optical linear diameters and the HI mass being non linear, the HI-deficiency parameter is not a distance independent value: for a given galaxy the HI-deficiency increases if its distance decreases."," The relationship between optical linear diameters and the HI mass being non linear, the HI-deficiency parameter is not a distance independent value: for a given galaxy the HI-deficiency increases if its distance decreases." + Garcia-Barreto et al. (, Garcia-Barreto et al. ( +"1994) used a distance of 10 Mpc in the determination of the HI mass of their sample, while we used the Tully-Fisher distance whenever available, or 14.52 Mpc elsewhere.","1994) used a distance of 10 Mpc in the determination of the HI mass of their sample, while we used the Tully-Fisher distance whenever available, or 14.52 Mpc elsewhere." + This difference in distance leads to an overestimate of the HI-deficiency parameter of ~ 0.04 for a typical Sc galaxy in the Garcia-Barreto et al., This difference in distance leads to an overestimate of the HI-deficiency parameter of $\sim$ 0.04 for a typical Sc galaxy in the Garcia-Barreto et al. + calculations with respect to ours., calculations with respect to ours. +" Conversely, the use of the calibration of Solanes et al. ("," Conversely, the use of the calibration of Solanes et al. (" +"1996) for Sa-Sc galaxies, which is based on Ho =100 km s! Mpc!, induces a decrease of the HI-deficiency parameter by a factor (1-d)Logh? (from 0.11 for Sa to 0.04 for Sc).","1996) for Sa-Sc galaxies, which is based on $H_0$ =100 km $^{-1}$ $^{-1}$, induces a decrease of the HI-deficiency parameter by a factor $d$ $h^2$ (from 0.11 for Sa to 0.04 for Sc)." +" Since the present sample is dominated by galaxies of type > Scd (78%)), whose distance has been determined using Ho =73 km s! Μρο”], the average HI—def is only marginally affected by the choice of Hy =100 km s! Mpc™! for Sa-Sc galaxies of Solanes et al. ("," Since the present sample is dominated by galaxies of type $\geq$ Scd ), whose distance has been determined using $H_0$ =73 km $^{-1}$ $^{-1}$, the average $HI-def$ is only marginally affected by the choice of $H_0$ =100 km $^{-1}$ $^{-1}$ for Sa-Sc galaxies of Solanes et al. (" +1996).,1996). +" The rest of the difference (0.18 in HI— def) might be due to statistical reasons, our sample (55 objects) being more than doubled with respect to that of Garcia-Barreto et al. ("," The rest of the difference (0.18 in $HI-def$ ) might be due to statistical reasons, our sample (55 objects) being more than doubled with respect to that of Garcia-Barreto et al. (" +"1994) (23 objects), or to the adopted calibration.","1994) (23 objects), or to the adopted calibration." + Garcia-Barreto et al. (, Garcia-Barreto et al. ( +1994) determined the HI-deficiency parameter using the B band luminosity relation of Giovanelli et al. (,1994) determined the HI-deficiency parameter using the B band luminosity relation of Giovanelli et al. ( +"1981), while our estimate is based on a diameter relation.","1981), while our estimate is based on a diameter relation." + The calibration of the HI-deficiency on optical diameters is less dispersed than that based on optical luminosities (Haynes Giovanelli 1984)., The calibration of the HI-deficiency on optical diameters is less dispersed than that based on optical luminosities (Haynes Giovanelli 1984). + We can thus conclude that late-type galaxies in the Coma I cloud are not as deficient in HI gas as previously claimed., We can thus conclude that late-type galaxies in the Coma I cloud are not as deficient in HI gas as previously claimed. + The Coma I cloud is thus composed of galaxies with a similar spiral fraction but richer in gas content than the Virgo M and W clouds., The Coma I cloud is thus composed of galaxies with a similar spiral fraction but richer in gas content than the Virgo M and W clouds. +" Being at a distance along the line of sight similar to that of Virgo (14.52 Mpc for Coma I and 16.5 Mpc for Virgo), and at a distance of ~ 5 Mpc on the plane of the sky to the core of the cluster, it could be considered as a cloud of Virgo (for comparison the M and W clouds are located at ~ 16 Mpc from the core of Virgo, Gavazzi et al."," Being at a distance along the line of sight similar to that of Virgo (14.52 Mpc for Coma I and 16.5 Mpc for Virgo), and at a distance of $\sim$ 5 Mpc on the plane of the sky to the core of the cluster, it could be considered as a cloud of Virgo (for comparison the M and W clouds are located at $\sim$ 16 Mpc from the core of Virgo, Gavazzi et al." + 1999)., 1999). + Is pre-processing, Is pre-processing +signature of a high accretion rate may appear as Fe K absorption in a highly ionised and massive outflow.,signature of a high accretion rate may appear as Fe K absorption in a highly ionised and massive outflow. + Hitherto most evidence for extreme outflows in AGN has been found in UV sudies of Broad Absorption Line (BAL) QSOs., Hitherto most evidence for extreme outflows in AGN has been found in UV sudies of Broad Absorption Line (BAL) QSOs. +" The observation reported here, of a massive high velocity outflow from bbroadens the scope of such studies."," The observation reported here, of a massive high velocity outflow from broadens the scope of such studies." +" BAL QSOs show absorption in a variety of, mainly high-ionisation, UV resonance transitions with velocity widths up to ~30000 km s! (e.g. Weymann 11991)."," BAL QSOs show absorption in a variety of, mainly high-ionisation, UV resonance transitions with velocity widths up to $\sim 30000$ km $^{-1}$ (e.g. Weymann 1991)." + About of optically selected QSOs display BALs., About of optically selected QSOs display BALs. +" As BAL QSOs appear otherwise similar to non-BAL QSOs, an ‘orientation model’ is traditionally invoked in which BAL QSOs are those in which the particular line-of-sight intersects an outflow which may be intrinsic to all QSOs."," As BAL QSOs appear otherwise similar to non-BAL QSOs, an `orientation model' is traditionally invoked in which BAL QSOs are those in which the particular line-of-sight intersects an outflow which may be intrinsic to all QSOs." + This model has recently been questioned by the discovery of a relatively high fraction (15 — 20%)) of radio-loud BAL quasars in the VLA FIRST survey bright quasar survey (Becker et al., This model has recently been questioned by the discovery of a relatively high fraction (15 – ) of radio-loud BAL quasars in the VLA FIRST survey bright quasar survey (Becker et al. + 2000)., 2000). + Becker et al., Becker et al. + propose BAL objects may be young or have recently been fuelled., propose BAL objects may be young or have recently been fuelled. +" In any case, the higher fraction of quasars that have BALs implies a higher fraction of the line-of-sight to the nucleus is covered with substantial absorbing material."," In any case, the higher fraction of quasars that have BALs implies a higher fraction of the line-of-sight to the nucleus is covered with substantial absorbing material." + Determining the amount of gas along the line-of-sight to a BAL is generally problematic due to a poor understanding of the relation between UV and X-ray absorption and the geometry of the flow., Determining the amount of gas along the line-of-sight to a BAL is generally problematic due to a poor understanding of the relation between UV and X-ray absorption and the geometry of the flow. +" Fitting UV absorption lines suggests Nu>107? cm~?, whereas the generally weak X-ray fluxes imply columns an order of magnitude or more higher (e.g. Hamann 1998; Sabra Hamann 2001; Gallagher 22002)."," Fitting UV absorption lines suggests $N_H \ge +10^{22}$ $^{-2}$, whereas the generally weak X-ray fluxes imply columns an order of magnitude or more higher (e.g. Hamann 1998; Sabra Hamann 2001; Gallagher 2002)." + Models in which the BAL gas is launched more or less vertically off a disk and then accelerated by radiation pressure (Murray 11995; Proga 22000) are reasonably consistent with the UV data but have difficulty in accelerating the large columns of material seen in X-rays unless they are launched from very close to the black hole - as we propose forPG1211--143.., Models in which the BAL gas is launched more or less vertically off a disk and then accelerated by radiation pressure (Murray 1995; Proga 2000) are reasonably consistent with the UV data but have difficulty in accelerating the large columns of material seen in X-rays unless they are launched from very close to the black hole - as we propose for. +" In summary, while PG1211--143 has strong soft X-ray emission and is not a BAL QSO in the UV, it does display a fast moving outflow and a line-of-sight column density which are similar to those required to explain, respectively, the UV and X-ray properties of BAL QSOs."," In summary, while PG1211+143 has strong soft X-ray emission and is not a BAL QSO in the UV, it does display a fast moving outflow and a line-of-sight column density which are similar to those required to explain, respectively, the UV and X-ray properties of BAL QSOs." +" Whether the outflow in bbecomes capable of producing BAL features further out in the flow, but we simply do not intersect such a line of sight, is unclear."," Whether the outflow in becomes capable of producing BAL features further out in the flow, but we simply do not intersect such a line of sight, is unclear." +" Neither do we yet know how common are X-ray absorption features as reported here forPG1211+143,, nor whether the X-ray absorbing gas causing BAL QSOs to be ‘X-ray weak’ is in outflow."," Neither do we yet know how common are X-ray absorption features as reported here for, nor whether the X-ray absorbing gas causing BAL QSOs to be `X-ray weak' is in outflow." +" However, it seems likely that the BAL phenomena and the high velocity outflow in aare closely related. ("," However, it seems likely that the BAL phenomena and the high velocity outflow in are closely related. (" +"1) An oobservation of the bright quasar hhas revealed evidence of a high velocity ionised outflow, with a mass and kinetic energy comparable to the accretion mass and bolometric luminosity, respectively. (","1) An observation of the bright quasar has revealed evidence of a high velocity ionised outflow, with a mass and kinetic energy comparable to the accretion mass and bolometric luminosity, respectively. (" +"2) A further implication of the high observed column density is that the inner flow is likely to be optically thick, providing a natural explanation for the strong BBB and soft X-ray emission inPG1211--143.. (","2) A further implication of the high observed column density is that the inner flow is likely to be optically thick, providing a natural explanation for the strong BBB and soft X-ray emission in. (" +"3) An extreme relativistic Fe K emission line apparent in a simple power law fit to the data can, alternatively, be explained in terms of partial covering of the continuum source by overlying matter in a lower ionisation state. (","3) An extreme relativistic Fe K emission line apparent in a simple power law fit to the data can, alternatively, be explained in terms of partial covering of the continuum source by overlying matter in a lower ionisation state. (" +4) We suggest the above properties might be common in AGN accreting at or close to the Eddington limit.,4) We suggest the above properties might be common in AGN accreting at or close to the Eddington limit. +" The results reported here are based on observations obtained withXMM-Newton,, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA)."," The results reported here are based on observations obtained with, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA)." + The authors wish to thank the SOC and SSC teams for organising the oobservations and initial data reduction and the referee for a careful and constructive reading of the initial text., The authors wish to thank the SOC and SSC teams for organising the observations and initial data reduction and the referee for a careful and constructive reading of the initial text. + ARK gratefully acknowledges a Royal Society Wolfson Research Merit Award., ARK gratefully acknowledges a Royal Society Wolfson Research Merit Award. +Dust huuimosities (Li;=Lis1000;12]) aud other piranueters are estimated using (Nolletal.2009))2.,"Dust luminosities $L_{\rm IR}=L [8{-}1000\,\mu{\rm m]}$ ) and other parameters are estimated using \citealt{noll09}) ." +. performs ai Bayesian analysis to estimate parameter by fitting models to the UV-to-subuuu SEDs., performs a Bayesian analysis to estimate parameter by fitting models to the UV-to-submm SEDs. + Oue can select amoug two sinele stellar sopulation librarics and several IR models/templates., One can select among two single stellar population libraries and several IR models/templates. + An AGN component can also be added to estimate he ACN fraction (contribution by a potential ACN to Lig), An AGN component can also be added to estimate the AGN fraction (contribution by a potential AGN to $L_{\rm IR}$ ). + The parameters of the dust attenuation law can )o mnodified aud allows for two separate stellar »»pulatious with a imultiphase dust treatiuent., The parameters of the dust attenuation law can be modified and allows for two separate stellar populations with a multiphase dust treatment. + We use the VLA GGIIz radio data from Morrisonal.(2010) which provides data down to a RAIS roise of 3.9 μὴν per beam., We use the VLA GHz radio data from \cite{morrison10} which provides data down to a RMS noise of 3.9 $\mu$ Jy per beam. +" The stellar ciunissiou is based on Maraston(2005).. while the dust emission is based ou Dale&Helou(2002) templates,"," The stellar emission is based on \cite{maraston05}, while the dust emission is based on \cite{dale02} templates." + provides dust ΠαππιοΊος Ly. while FUV huuinosities Lppis are derived at Ager= 1521120 and are defined through the quantity νο.," provides dust luminosities $L_{\rm IR}$, while FUV luminosities $L_{\rm FUV}$ are derived at $\lambda_{\rm rest}=153\,$ nm and are defined through the quantity $\nu L_{\nu}$ ." + An cucrev budget is performed οιιο the ft. and the πιαπι value allowed for Lj has to be consistent with the energv moved by dust erains from the UV-optical range to the FIR raus.," An energy budget is performed during the fit, and the maximum value allowed for $L_{\rm IR}$ has to be consistent with the energy moved by dust grains from the UV-optical range to the FIR range." + Tudividual SEDs with the best models selected by are shown in Fie., Individual SEDs with the best models selected by are shown in Fig. + 35. for the five LBGs with radio data and for the (:= 1.9) LBG., \ref{sedlbg} for the five LBGs with radio data and for the $z=1.9$ ) LBG. + Observational and plysical parameters are given in Table 1.., Observational and physical parameters are given in Table \ref{Table}. + Fig., Fig. + 3 shows that is able to fit the SEDs from the FUV o the radio successfully., \ref{sedlbg} shows that is able to fit the SEDs from the FUV to the radio successfully. + The average PUY luminosity of the SPIRE detected objects is lost(Lgvv/b.);=(10.7+ 0.2). and most of them are therefore UV- Calaxics.," The average FUV luminosity of the SPIRE detected objects is $ \log\langle(L_{\rm FUV}/{\rm L}_\odot)\rangle = (10.7 \pm 0.2)$ , and most of them are therefore UV-Luminous Galaxies." + Their average dust luminosity is ogtCLis/E.)=(11.9τε} aud their stellar average uass is logétM./Mi)=(GHOd 05)., Their average dust luminosity is $ \log\langle(L_{\rm IR}/{\rm L}_\odot) \rangle= (11.9 \pm 0.1)$ and their stellar average mass is $ \log\langle(M_*/{\rm M}_\odot) \rangle = (11.0 \pm 0.5)$ . + A comparison with Maedisetal.(2010). shows that the average stellar 11ass is similar to the average stellar iuass (log(AL/M9)= 11) of :~3 LDCGs detected at A= San. The dust temperatures Ty are estimated for a few objects by fitting nodified black bodies with au eumidssivitv iudex of 1.5.," A comparison with \cite{magdis10} shows that the average stellar mass is similar to the average stellar mass $ \log\langle(M_*/{\rm M}_\odot) \rangle = 11$ ) of $z \sim 3$ LBGs detected at $\lambda = 8\,\mu$ m. The dust temperatures $T_{\rm d}$ are estimated for a few objects by fitting modified black bodies with an emissivity index of 1.5." + We find that two of the low-: LBCs have Z4 ~ 211s. while he hieh-: ULIRG has Z4= 52K (see Table 11).," We find that two of the $z$ LBGs have $T_{\rm d}$ $\sim24\,$ K, while the $z$ ULIRG has $T_{\rm d} = 52\,$ K (see Table \ref{Table}) )." + The wo low-: LBGCs have quite low temperatures compared with ULIRGS/SAIGs. but may be typical of star-forming ealaxies in eeuceral. while the high-+ oue is similar to e.g. Arp220.," The two $z$ LBGs have quite low temperatures compared with ULIRGS/SMGs, but may be typical of star-forming galaxies in general, while the $z$ one is similar to e.g. Arp220." +" Ow LBC SEDs in the rest-frame UV are dn very eood aegrecincnut with the rest-frame UV spectra of the Shapleyetal. (2003)""s composite spectrum at i72."," Our LBG SEDs in the rest-frame UV are in very good agreement with the rest-frame UV spectra of the \cite{shapley03}' 's composite spectrum at $z\,{\sim}\,3$." + Whatever the redshift range. LBCs secur to preseut the same starburst characteristics in the rest-frame UV. which is expected. given that similar rest-frame color selection criteria are used.," Whatever the redshift range, LBGs seem to present the same starburst characteristics in the rest-frame UV, which is expected, given that similar rest-frame color selection criteria are used." + However. the colors become redder when the objects are detected at 1402 by and even redder when they are detected at yan bv SPIRE.," However, the colors become redder when the objects are detected at $\,\mu$ m by and even redder when they are detected at $\,\mu$ m by SPIRE." + This cau be interpreted as being due to ligher dust attenuations (see Fie. 23).," This can be interpreted as being due to higher dust attenuations (see Fig. \ref{color}) )," + aud is consistent with Burearellaetal.(2007).. who fouud that Spitzer-detected and undetected LBCs have about the same stellar population ages. but the latter are more extinguished.," and is consistent with \cite{burgarella07}, who found that -detected and undetected LBGs have about the same stellar population ages, but the latter are more extinguished." + The Δέν fraction determined by is always consistent withzero. except for the galaxy | 62083L3. for which a ACN coutribution to the FIR dust huninosity is sugeested.," The AGN fraction determined by is always consistent withzero, except for the galaxy $+$ 620834.3, for which a AGN contribution to the FIR dust luminosity is suggested." + Most of the IR excess due to a potential AGN should be in the MIR. while the SPIRE flux is expected to be dominated by the starburst component (Ilatziuunaoeglouetal.2010):: accounts for both.," Most of the IR excess due to a potential AGN should be in the MIR, while the SPIRE flux is expected to be dominated by the starburst component \citep{hatziminaoglou10}; accounts for both." + The vast majority of the dropout galaxies are not detected by SPIRE with the preseut detection limits., The vast majority of the dropout galaxies are not detected by SPIRE with the present detection limits. + If LBCs were to follow Meurer.Beckman.&Calzetti(1999) relation (rolatiug UV-attenuation to dust cuiission). what fraction of them would be detectable?," If LBGs were to follow \cite{meurer99} relation (relating UV-attenuation to dust emission), what fraction of them would be detectable?" + We can use the (VW) color as a proxy for the rest-frame m{PUV]ΜΗΝ θαΊος, We can use the $(U - V)$ color as a proxy for the rest-frame $m[{\rm FUV}] - m[{\rm NUV}]$. + suggest that this color provides > with an accuracy better than 0.01 for power-laws fyxwAT and }=2. Land 0.," Simulations suggest that this color provides $\beta$ with an accuracy better than 0.01 for power-laws $f_\lambda \propto \lambda^\beta$ and $\beta = -2$, $-1$ and 0." + In detail. we can estimate: From the > values. we estimate loe(Lig/Lpry) for the LBCs and. after evaluating Levy from the filter closest to Apry=O.15 umm (as a function of the redshift}. we determine Ly.," In detail, we can estimate: From the $\beta$ values, we estimate $\log (L_{\rm IR}/L_{\rm FUV}$ ) for the LBGs and, after evaluating $L_{\rm FUV}$ from the filter closest to $\lambda_{\rm FUV}$ nm (as a function of the redshift), we determine $L_{\rm IR}$ ." + These values of Lig can be transformed into Losy and Sosy usine the following calibrations coluputed from the Dale&Helou(2002) models with LO0 will be considered. later."," First, ${\cal N}_m(\xi)$ is an even function of $\xi$ so only $\xi\ge 0$ will be considered later." +"Secondly. A,(£)decreases monotonically with increasing £>»QO.","Secondly, ${\cal N}_m(\xi)$decreases monotonically with increasing $\xi>0$." +" Thirdly. 0«NS,<1 for m>] while No is positive and can be >greater than | for a sullicientlv small £."," Thirdly, $0<{\cal N}_m<1$ for $m\ge 1$ while ${\cal N}_0$ is positive and can be greater than 1 for a sufficiently small $\xi$." + The choice of such. form. of perturbations is. differentfrom that of νο ‘Tremaine (1996). whose spiral perturbations were taken to be pxr1o24exptim£-—Inr) [or mz»0 (analysis before subsection 3.4 in their paper) in our notations.," The choice of such form of perturbations is differentfrom that of Syer Tremaine (1996), whose spiral perturbations were taken to be $\mu\propto r^{-1-2\beta}\exp(\hbox{i}m\xi\ln r)$ for $m>0$ (analysis before subsection 3.4 in their paper) in our notations." + For axisvmmetrie stability analysis in their subsection 3.4. they adopted the same spiral perturbations in the form of (62)) (see also Lemos et al.," For axisymmetric stability analysis in their subsection 3.4, they adopted the same spiral perturbations in the form of \ref{spirPerturb}) ) (see also Lemos et al." + 1991)., 1991). + We note hat our background equilibria as well as the adopted form of logarithmic spiral perturbations are themselves. scale-ree. separately. whereas combinations of the background equilibrium and. perturbations are not scale-free except. for he special 7=1/4case (see Lynden-Bell Lemos 1993).," We note that our background equilibria as well as the adopted form of logarithmic spiral perturbations are themselves scale-free, separately, whereas combinations of the background equilibrium and perturbations are not scale-free except for the special $\beta=1/4$case (see Lynden-Bell Lemos 1993)." + Parallelling with ο for the case of aligned yverturbations. there are two useful formulae. Lor AN(E) or logarithmic spiral perturbations.," Parallelling with ${\cal P}_m$ for the case of aligned perturbations, there are two useful formulae for ${\cal N}_m(\xi)$ for logarithmic spiral perturbations." + The first one is the recursion relation [or m?|€5»L1., The first one is the recursion relation and the second one is the asymptotic expression for ${\cal N}_m(\xi)$ for $m^2+\xi^2\gg 1$. +" For m©2. this asvmptotic expression (65)) is accurate enough to compute values of AS,(E)."," For $m\ge 2$, this asymptotic expression \ref{asymN}) ) is accurate enough to compute values of ${\cal N}_m(\xi)$." + Using potential-densitv set. of (62)) for. logarithmic spirals. we rearrange stationary coplanar perturbation equations (27))and (28)) into the following forms with mc.," Using potential-density set of \ref{spirPerturb}) ) for logarithmic spirals, we rearrange stationary coplanar perturbation equations \ref{stationarys}) )and \ref{stationaryg}) ) into the following forms with $m>0$." + where the four relevant coellicients 7/4. 445. Gy and Cs are defined here explicitly by Rewriting equations (66)) with expressions (62)) for 40 and pe. we immediately. obtain As for the aligned case in subsection 3.1. we define some usefulnotations for parameter combinations that will simplify our following derivations. namely With convenient.notations (69)). equation (68)) leads to the following stationary clispersion relation for coplanar logarithmic spiral perturbations in a composite disc system.," where the four relevant coefficients $H_1$, $H_2$, $G_1$ and $G_2$ are defined here explicitly by Rewriting equations \ref{hoho3}) ) with expressions \ref{spirPerturb}) ) for $\mu^s$ and $\mu^g$, we immediately obtain As for the aligned case in subsection 3.1, we define some usefulnotations for parameter combinations that will simplify our following derivations, namely With convenientnotations \ref{ABCHspiral}) ), equation \ref{hoho4}) ) leads to the following stationary dispersion relation for coplanar logarithmic spiral perturbations in a composite disc system." +" Substituting expressions (67)) of ἐν. H5. €, and C» into stationary dispersion. relation (70)) ancl using the background condition DF=g(D;|1). I. we obtain one quacratic equation in terms of jy—Dz. namely where coellicients C5. Cy, and Cy are functions of parameters m. 4 2. ὃν g and £. and are defined by Given specific values lor a. S. 0 and η. we readily solve quacdratie equation (71)) analytically for cach ""racial wavenumber £ forstationary logarithmic spirals."," Substituting expressions \ref{spiralHG}) ) of $H_1$, $H_2$, $G_1$ and $G_2$ into stationary dispersion relation \ref{SDPspiral}) ) and using the background condition $D_g^2=\eta(D_s^2+1)-1$ , we obtain one quadratic equation in terms of $y\equiv D_s^2$, namely where coefficients $C_2$ , $C_1$ and $C_0$ are functions of parameters $m$ , $\beta$ , $\delta$ , $\eta$ and $\xi$ , and are defined by Given specific values for $m$ , $\beta$ , $\delta$ and $\eta$ , we readily solve quadratic equation \ref{spiral}) ) analytically for each `radial wavenumber' $\xi$ forstationary logarithmic spirals." + Again for a non-negative determinant as proven in Appendix A. there," Again for a non-negative determinant as proven in Appendix A, there" +produces a stuall component of very hot eas surrounding the stellar population. so hot that its cooling time is very long.,"produces a small component of very hot gas surrounding the stellar population, so hot that its cooling time is very long." + Untortunately at the resolition currently achievable in cosmological simulations. this primitive thermal prescription for feedback deposits the sauce SN enerev iuto a immuch larger reservoir of eas. which does not reach the same high temperature as it should.," Unfortunately at the resolution currently achievable in cosmological simulations, this primitive thermal prescription for feedback deposits the same SN energy into a much larger reservoir of gas, which does not reach the same high temperature as it should." +" This now-warmn componcut of eas can casily radiate the excess enerev away. cool further aud proceed with star-formation runaway thus defeating the purpose of the feedback (οον,2).."," This now-warm component of gas can easily radiate the excess energy away, cool further and proceed with star-formation runaway thus defeating the purpose of the feedback \citep[e.g.,][]{Steinmetz:1999p635}." + Building from these failures. a nunber of research eroups artificially turn off cooling iu a gas parcel for a period of time (f.~10* vr) after a cluster of stars has formed out of it (c.g...77????)..," Building from these failures, a number of research groups artificially turn off cooling in a gas parcel for a period of time $t \sim 10^7$ yr) after a cluster of stars has formed out of it \citep[e.g.,][]{Gerritsen:1997p1039, Thacker:2000p1040, SommerLarsen:2003p1116, Stinson:2006p1023, Governato:2007p1022, Agertz:2010p461, Colin:2010p1053, Piontek:2011p1041, Guedes:2011p1080}." +" Phas method is justified as an application of the77? Sedov-Tavlor blast wave solution for a Type II SN δν, which blows out any cold media from the imuuecdiate euvironnment of a star formation event."," This method is justified as an application of the Sedov-Taylor blast wave solution for a Type II SN \citep{Taylor:1950p1075, Sedov:1959p1074}, which blows out any cold media from the immediate environment of a star formation event." + Using this prescription. any eas 1u a galaxy which starts to collapse iuto knots will reach he star formation criteria. form a star. and then heat up without auv allowed cooling. thus preventing further collapse.," Using this prescription, any gas in a galaxy which starts to collapse into knots will reach the star formation criteria, form a star, and then heat up without any allowed cooling, thus preventing further collapse." + Not surprisingly. these research groups have ound some success With this method. viclding simulated ealaxies with reduced πι rotation curves due to less nassive bulee conipoueuts: however. gas parcel niasses and sizes in cosmological sinulatious of this sort are vpically too lavee for the Sedov-Tavlor solution to apply (see Section 2.1.0)).," Not surprisingly, these research groups have found some success with this method, yielding simulated galaxies with reduced inner rotation curves due to less massive bulge components; however, gas parcel masses and sizes in cosmological simulations of this sort are typically too large for the Sedov-Taylor solution to apply (see Section \ref{code_coolingsuppression}) )." + Thus despite the successes of he cooling suppression feedback model. the σοιτν continues to search for other more plivsicallv-1uotivated solutious.," Thus despite the successes of the cooling suppression feedback model, the community continues to search for other more physically-motivated solutions." + Another suberid imodel for feedback (i.e. on scales stnaller than the true resolution of the simulation) is to inject kinetic cuerey directly into the eas: this cau alleviate the problem of thermal cucrey beiug radiated away., Another subgrid model for feedback (i.e. on scales smaller than the true resolution of the simulation) is to inject kinetic energy directly into the gas; this can alleviate the problem of thermal energy being radiated away. + For example. some studies (c.g.77?7) MSIE Smoothed Particle Ibdrodyvuaiics (SPIT) give some of the SN enerey to individual eas particles iu the form of momentum.," For example, some studies \citep[e.g.][]{Springel:2003p1044, Scannapieco:2006p1118, Oppenheimer:2008bu} using Smoothed Particle Hydrodynamics (SPH) give some of the SN energy to individual gas particles in the form of momentum." + This method cau result in significant mass outflows (bv desigu). but at the cost of decoupling wiud particles from livcodvuamic interaction for a period of time.," This method can result in significant mass outflows (by design), but at the cost of decoupling wind particles from hydrodynamic interaction for a period of time." + An alternate approach. to keep wind particles coupled to the disk gas was explored by ?..," An alternate approach, to keep wind particles coupled to the disk gas was explored by \citet{Schaye:2008p1045}." + Both approaches help but. by themselves. do not appear to eoncrate realistic rotation curves.," Both approaches help but, by themselves, do not appear to generate realistic rotation curves." + Iu addition. ? showed that insufficient resolution in a simulation can lead to artificial fragmentation of the gas. perhaps resulting ina further overproduction of stars.," In addition, \citet{Truelove:1997p1046} showed that insufficient resolution in a simulation can lead to artificial fragmentation of the gas, perhaps resulting in a further overproduction of stars." + One wav to prevent artificial fragmentation is to add additional (ummerical) pressure in lieh-deusitv. low-telmperature regions to ensure that the Jeaus leneth is alwavs resolved (2)...," One way to prevent artificial fragmentation is to add additional (numerical) pressure in high-density, low-temperature regions to ensure that the Jeans length is always resolved \citep{Machacek:2001p1047, Robertson:2008p1017}." + This can be achieved by modifving the equation of state (EOS) itself. maline it stiffer iu order to provide au additional source of pressure to gas in deuser reeious (227)...," This can be achieved by modifying the equation of state (EOS) itself, making it stiffer in order to provide an additional source of pressure to gas in denser regions \citep{Schaye:2008p1045, Ceverino:2009p1014, Agertz:2010p461}." + A polytropic EOS (P.x pF) with P=1/3 will keep the ratio of Jeaus leneth to resolution leugth constant (assunius Laeraneian resolution such that the resolution leusth decreases as p17 ων fixe1 resolution D—2 is required). but even stiffer relations have been used.," A polytropic EOS $P \propto \rho^{\Gamma}$ ) with $\Gamma = 4/3$ will keep the ratio of Jeans length to resolution length constant (assuming Lagrangian resolution such that the resolution length decreases as $\rho^{-1/3}$ – for fixed resolution $\Gamma = 2$ is required), but even stiffer relations have been used." + For example. ? rau simulations wit1 such an equation of state. where in low-density regions it behaved as an ideal gas. but in high-density (star forming) regions it followed a polytropic equation of state with P2.," For example, \citet{Agertz:2010p461} ran simulations with such an equation of state, where in low-density regions it behaved as an ideal gas, but in high-density (star forming) regions it followed a polytropic equation of state with $\Gamma = 2$." + Iu this paper. we undertake an investigation of galaxy formation using an Adaptive Mesh. Refinement (AMR) lydrodvuamics code.," In this paper, we undertake an investigation of galaxy formation using an Adaptive Mesh Refinement (AMR) hydrodynamics code." + The majority of work in this field has used SPIT codes. aud so this allows us to investigate the problem from a new augle.," The majority of work in this field has used SPH codes, and so this allows us to investigate the problem from a new angle." + Although there has been some work with AMIR codes (2727).. there has not been a clear demonstration that an equivalent AAIR calculation (i.c. one without a suberid feedback model) actually docs reproduce the classic SPIT result.," Although there has been some work with AMR codes \citep{Joung:2009p1010, Ceverino:2009p1014, Agertz:2010p461, Colin:2010p1053}, there has not been a clear demonstration that an equivalent AMR calculation (i.e. one without a subgrid feedback model) actually does reproduce the classic SPH result." + We begin by simulating a set of five halos without any ecdhack or suberid model (except a muni pressure support to prevent artificial fragmentation)., We begin by simulating a set of five halos without any feedback or subgrid model (except a minimum pressure support to prevent artificial fragmentation). + We find. in agreement with SPILT codes that a large. concentrated πι]οο is produced. resulting iu a rotation curve that rises to ~500 kin/s at r~1l kpe.," We find, in agreement with SPH codes that a large, concentrated bulge is produced, resulting in a rotation curve that rises to $\sim 500$ km/s at $r \sim 1$ kpc." + We then vary a uber of nunucerical aud plysical parameters iu order o understand how sensitive the result is to our a choice of parameters., We then vary a number of numerical and physical parameters in order to understand how sensitive the result is to our a choice of parameters. + The paper is organized as follows., The paper is organized as follows. + Section 2. describes he details of our lvdrodvuamics code. our initial conditious aud the relevant parameters for this study.," Section \ref{method} describes the details of our hydrodynamics code, our initial conditions and the relevant parameters for this study." + Iu Section 3.. we present the results of our simulations including the five canonical runs. our resolution study. and our imnodiB&Bed runs.," In Section \ref{results}, we present the results of our simulations including the five canonical runs, our resolution study, and our modified runs." + Section [. ds a discussion of our results and their duplications., Section \ref{discussion} is a discussion of our results and their implications. + Finally. Section 5 sunuuamrizes our couclusions aud makes predictions for future solutions to the aneular momentum problem.," Finally, Section \ref{summary} summarizes our conclusions and makes predictions for future solutions to the angular momentum problem." + Our simulations were performed usingfootuoteenzo-project.ore.. anu Eulerian. threc-dineusional. exid-based lvdrodvuamics code that cluplovs adaptive mesh refinement— iu order to achieve/ targeted regious of high resolution iu a cosmological volume (77)..," Our simulations were performed using, an Eulerian, three-dimensional, grid-based hydrodynamics code that employs adaptive mesh refinement in order to achieve targeted regions of high resolution in a cosmological volume \citep{Bryan:1997p869, OShea:2004p446}." + Gas is) discretized on the eric. but dark iatter and stars are treated as particles.," Gas is discretized on the grid, but dark matter and stars are treated as particles." + The lvdrodvuamics code (7) is used to evolve the eas on the erid., The hydrodynamics code \citep{Stone:1992p1117} is used to evolve the gas on the grid. + inclides eas. sclferavity. a non-equilibrium model for IT aud Πο ionization and cooling. a metagalactic ultraviolet backerouud (7).. aud equilibriun cooling due to metals (although for the runs described in this paper. we do not include metal cooling).," includes gas, self-gravity, a non-equilibrium model for H and He ionization and cooling, a metagalactic ultraviolet background \citep{Haardt:1996p1000}, and equilibrium cooling due to metals (although for the runs described in this paper, we do not include metal cooling)." + Star formation is modeled using a simple criteria based on ?.., Star formation is modeled using a simple criteria based on \citet{Cen:1992p1071}. + A exid cell will produce a star if (1) the overdensity in that cell exceeds a given value (og). (ii) the mass of eas in the cell exceeds the local Jeans mass. (d) there is locally convergent flow (ie. the velocity divergence is negativo) and (v) the cooling time for the gas to collapse is less than the dynamical time in that cell (or the temperature is near the nuünimuun allowed. arouud 104 Ky.," A grid cell will produce a star if: (i) the overdensity in that cell exceeds a given value $\delta_{\rm SF}$ ), (ii) the mass of gas in the cell exceeds the local Jeans mass, (iii) there is locally convergent flow (i.e. the velocity divergence is negative) and (iv) the cooling time for the gas to collapse is less than the dynamical time in that cell (or the temperature is near the minimum allowed, around $10^4$ K)." +" HE a exid cell meets all the previous criteria then eas is converted ito a ""star particle”. as calculated using where egg is the star formation cficiency (nore properly the efficieucy per dynamical time). At is the size"," If a grid cell meets all the previous criteria then gas is converted into a “star particle”, as calculated using where $\epsilon_{\rm SF}$ is the star formation efficiency (more properly the efficiency per dynamical time), $\Delta t$ is the size" +homogeneous. foreground. svuchrotron emission region respousible for the IR flare.,"homogeneous, foreground, synchrotron emission region responsible for the IR flare." + We characterize, We characterize +the parameters adopted (see Fig. 2)).,the parameters adopted (see Fig. \ref{fig_comp}) ). +" The internal energy Uxn,. and the cooling rate Fn2. which indicates that the cooling timescale increases with decreasing electron number density n.."," The internal energy $U\propto +n_{\rm e}$, and the cooling rate $F^{-}\propto n_{\rm e}^2$, which indicates that the cooling timescale increases with decreasing electron number density $n_{\rm e}$." + In the estimate of the cooling. we assume that the radial velocity of the outflow is the same as the virialized velocity. which is the least velocity that the outflow can escape to infinity.," In the estimate of the cooling, we assume that the radial velocity of the outflow is the same as the virialized velocity, which is the least velocity that the outflow can escape to infinity." +" If the gases in the outflow move at the speed higher than the virialized. velocity. the number density 7, of the electrons decreases with increasing outflow velocity provided all other parameters are fixed. and therefore the cooling timescale becomes larger for higher outflow velocity."," If the gases in the outflow move at the speed higher than the virialized velocity, the number density $n_{\rm e}$ of the electrons decreases with increasing outflow velocity provided all other parameters are fixed, and therefore the cooling timescale becomes larger for higher outflow velocity." + The results plotted in Fig., The results plotted in Fig. +" | are caleulated with jj,=1. Le. M,=M. and m=imu0.01. which. of course. leads to an lower limit on the cooling length scale (see Eqs."," \ref{fig_r_blr} are calculated with $\eta_{\rm w}=1$, i.e., $\dot{M}_{\rm w}=\dot{M}$, and $\dot{m}=\dot{m}_{\rm crit}=0.01$, which, of course, leads to an lower limit on the cooling length scale (see Eqs." + 12 and 14)).," \ref{l_cool_min} + and \ref{l_cool_min3}) )." + For most of the LLAGNs. the two parameters. jj«| and m«moa. are satisfied. which strengthens the conclusion derived m our estimates.," For most of the LLAGNs, the two parameters, $\eta_{\rm w}\ll 1$ and $\dot{m}\ll \dot{m}_{\rm crit}$, are satisfied, which strengthens the conclusion derived in our estimates." + The detailed physics for the transition of accretion modes is still unclear., The detailed physics for the transition of accretion modes is still unclear. + It was suggested that the ADAF co-exists with the standard disk. r.e.. the inner ADAF connects to the outer thin aceretion disk. in some sources accreting at rates slightly lower than the critical rate 7744 (e.g..Quataertetal.1999:Cao2003;Xu&Cao 2009)..," It was suggested that the ADAF co-exists with the standard disk, i.e., the inner ADAF connects to the outer thin accretion disk, in some sources accreting at rates slightly lower than the critical rate $\dot{m}_{\rm crit}$ \citep*[e.g.,][]{1999ApJ...525L..89Q,2003ApJ...599..147C,2009RAA.....9..401X}." + The transition radius increases with decreasing accretion rate i. which is expected by the thermal instability or disk evaporation induced transition scenarios (e.g..Abramowiczetal.1995:Liu1999:RózaNska&Czerny2000:SpruitDeufel 2002)..," The transition radius increases with decreasing accretion rate $\dot{m}$, which is expected by the thermal instability or disk evaporation induced transition scenarios \citep*[e.g.,][]{1995ApJ...438L..37A,1999ApJ...527L..17L,2000A&A...360.1170R,2002A&A...387..918S}." + In the presence of an outer cold disk. the soft photons from the cold disk will be Compton upseattered by the hot electrons in the outflow.," In the presence of an outer cold disk, the soft photons from the cold disk will be Compton upscattered by the hot electrons in the outflow." + For the cases that the radiatively& cooling can be neglected. the temperature of the gas will drop in an adiabatically expanding outflow.," For the cases that the radiatively cooling can be neglected, the temperature of the gas will drop in an adiabatically expanding outflow." +" Our estimate shows that the gas temperature 74,xrin the outflow."," Our estimate shows that the gas temperature $T_{\rm +gas}\propto r^{-1}$ in the outflow." +" The typical temperature of the ions in an ADAF near the black hole is ~I0! K (e.g.MeClintock2008).. the gases can be cooled to the typical temperature of BLRs (~10* K) only in the outflow with a distance >10°"" Schwarzschild radii from the black hole."," The typical temperature of the ions in an ADAF near the black hole is $\sim 10^{11-12}$ K \citep*[e.g.][]{2008NewAR..51..733N}, the gases can be cooled to the typical temperature of BLRs $\sim 10^4$ K) only in the outflow with a distance $>10^{6-7}$ Schwarzschild radii from the black hole." + It corresponds to ~1077 light days for a black hole with M=10'M... which is obviously beyond the BLR in luminous AGNs (see Fig. 1).," It corresponds to $\sim 10^{4-5}$ light days for a black hole with $M=10^{7}M_\odot$, which is obviously beyond the BLR in luminous AGNs (see Fig. \ref{fig_r_blr}) )." + Therefore. we propose that the outflows from the ADAFs in LLAGNSs are too hot to be cooled to form clouds in the BLRs. which leads to the disappearance of the BLR in LLAGNs.," Therefore, we propose that the outflows from the ADAFs in LLAGNs are too hot to be cooled to form clouds in the BLRs, which leads to the disappearance of the BLR in LLAGNs." + À small fraction of AGNs were found to have emission lines with double-peaked profiles (e.g..Eracleous&Halpern1994:Stratevaetal. 2003).. which usually have low Eddington ratios (seeEracleous2006.forareviewandreferencestherein.butalsoseeWu&Liu2004;Bianetal. 2007)..," A small fraction of AGNs were found to have emission lines with double-peaked profiles \citep*[e.g.,][]{1994ApJS...90....1E,2003AJ....126.1720S}, which usually have low Eddington ratios \citep*[see][for a review and +references therein, but also see Wu \& Liu 2004; Bian et al. +2007]{2006ASPC..360..217E}." + The most favorite model for the double-peaked emitters suggests that the double-peaked broad emission lines are emitted from a ring in the accretion disk. which may also be photo-1onized by the radiation from the inner region or/and the outflow (e.g..Chenetal.1959;Nemmen2006:Cao&Wang 2006)..," The most favorite model for the double-peaked emitters suggests that the double-peaked broad emission lines are emitted from a ring in the accretion disk, which may also be photo-ionized by the radiation from the inner region or/and the outflow \citep*[e.g.,][]{1989ApJ...339..742C,2006ApJ...643..652N,2006ApJ...652..112C}." + The observed broad-line emission may originate from two separated regions: the clouds in the normal BLRs. or/and the outer ring in the thin aceretion disk.," The observed broad-line emission may originate from two separated regions: the clouds in the normal BLRs, or/and the outer ring in the thin accretion disk." + The broad-line emission from the BLR clouds dominates over that from the outer region of the accretion disk in normal broad-line AGNs., The broad-line emission from the BLR clouds dominates over that from the outer region of the accretion disk in normal broad-line AGNs. +" For the double-peaked emitters accreting at rates lower than the critical accretion rate 71,4. the ADAF is present in the inner region and connects to the outer thin accretion disk."," For the double-peaked emitters accreting at rates lower than the critical accretion rate $\dot{m}_{\rm crit}$, the ADAF is present in the inner region and connects to the outer thin accretion disk." + The gases in the outflow from the ADAF are too hot to be cooled to form the clouds in the BLR when the transition radius of the ADAF to the outer disk ru220 and the secondary small cold disk is less luminous than Loy&0.0032 which leads to the disappearance of BLR clouds in these ug.sources.," The gases in the outflow from the ADAF are too hot to be cooled to form the clouds in the BLR when the transition radius of the ADAF to the outer disk $r_{\rm d,tr}\ga20$ and the secondary small cold disk is less luminous than $L_{\rm sd}\la +0.003L_{\rm Edd}$, which leads to the disappearance of BLR clouds in these sources." + Thus. the line emission from the outer region of the accretion disk is not contaminated by the emission from the BLR clouds. which emerges as double-peaked emission lines.," Thus, the line emission from the outer region of the accretion disk is not contaminated by the emission from the BLR clouds, which emerges as double-peaked emission lines." + This also provides a clue to the theoretical models for the accretion mode transition., This also provides a clue to the theoretical models for the accretion mode transition. + I thank the referee for the νου helpful comments/suggestions., I thank the referee for the very helpful comments/suggestions. + This work is supported by the— NSFC (grants 10773020. 10821302. and 10833002). the National Basie Research Program of China (grant 2009CB824800). the Science and Technology Commission of Shanghai Municipality (10XD1405000).. the CAS (grant KJCX2- and the CAS/SAFEA International Partnership Program for Creative Research Teams.," This work is supported by the NSFC (grants 10773020, 10821302, and 10833002), the National Basic Research Program of China (grant 2009CB824800), the Science and Technology Commission of Shanghai Municipality (10XD1405000), the CAS (grant KJCX2-YW-T03), and the CAS/SAFEA International Partnership Program for Creative Research Teams." +We have chosen spectra containing the interstellar molecular bands of Cs and the DIB at towards eleven objects acquired using different instruments.,We have chosen spectra containing the interstellar molecular bands of $C_2$ and the DIB at towards eleven objects acquired using different instruments. + Objects with intermediate colour excess and a single (or one strongly dominant) velocity component were selected (Table 1) to make it likely that analysed features originate from single clouds., Objects with intermediate colour excess and a single (or one strongly dominant) velocity component were selected (Table 1) to make it likely that analysed features originate from single clouds. + The profiles of K7. CH and CH lines were checked for the existence of more than one dominating Doppler component visible in our high resolution and high signal-to-noise ratio spectra.," The profiles of $KI$ , $CH$ and $CH^+$ lines were checked for the existence of more than one dominating Doppler component visible in our high resolution and high signal-to-noise ratio spectra." + Almost all of these lines in the spectra of the programme stars are free of any observable Doppler splitting., Almost all of these lines in the spectra of the programme stars are free of any observable Doppler splitting. + One exception is the 6Η line towards 2204827., One exception is the $CH^+$ line towards 204827. + Previous papers (??) show multiple components in very high resolution spectra in the interstellar lines towards 2204827 and a very weak Doppler component in KI and CH toward 1148184.," Previous papers \citep{Pan2004,Welty2001} show multiple components in very high resolution spectra in the interstellar lines towards 204827 and a very weak Doppler component in KI and CH toward 148184." + These authors also show that almost all of the program objects have multiple components in Nal: however such components may be unrelated to molecular ones or to DIBs (?).., These authors also show that almost all of the program objects have multiple components in NaI; however such components may be unrelated to molecular ones or to DIBs \citep{Bondar2007}. + Moreover. the latest survey of diffuse interstellar bands was based on 2204827 spectra (?)..," Moreover, the latest survey of diffuse interstellar bands was based on 204827 spectra \citep{Hobbs2008}." + The lack of Doppler splitting in the interstellar CH AA)) line in our spectra is demonstrated in Figure | (see also Doppler splitting is easier to detect when the observed feature is narrower., The lack of Doppler splitting in the interstellar $CH$ ) line in our spectra is demonstrated in Figure 1 (see also Doppler splitting is easier to detect when the observed feature is narrower. + Thus. while the observed CH profiles do not show Doppler components. the latter cannot cause the observed broadening of the much broader DIB.," Thus, while the observed CH profiles do not show Doppler components, the latter cannot cause the observed broadening of the much broader DIB." + Naturally it is not possible to prove that DIB carriers are spatially correlated to any atomic and molecular species: the only hints come from their similar Doppler shifts in different objects., Naturally it is not possible to prove that DIB carriers are spatially correlated to any atomic and molecular species; the only hints come from their similar Doppler shifts in different objects. + This paper is based on six sets of observedspectra., This paper is based on six sets of observedspectra. + Strengths of the Phillips ») bands of the C» molecule were measured in the spectra from UVES (?). (one exception is 223180 where we used data from BOES). because this spectrograph allows us to observe three vibrational bands of this system.," Strengths of the Phillips ) bands of the $C_2$ molecule were measured in the spectra from UVES \citep{Bagnulo2003} (one exception is 23180 where we used data from BOES), because this spectrograph allows us to observe three vibrational bands of this system." + We measured the absorption lines from the three bands (where possible) of C» (1-0)A.. (2-0) and (3-0) to make our estimates of the excitation temperature more reliable.," We measured the absorption lines from the three bands (where possible) of $C_2$ (1-0), (2-0) and (3-0) to make our estimates of the excitation temperature more reliable." + For 1110432. 1147888. 1149757 and 2204827 the excitation temperature of C» was taken from previous papers (from the Phillips system as well as the Mulliken band: see Table |).," For 110432, 147888, 149757 and 204827 the excitation temperature of $C_2$ was taken from previous papers (from the Phillips system as well as the Mulliken band; see Table 1)." + The diffuse interstellar band at was analysed using the spectra from GECKO. FEROS. HARPS. CES. UVES and BOES depending on which one was available or had a better signal-to-noise ratio in this range of the spectrum. as shown in One of the set of spectra used to analyse the interstellar absorption. features was acquired with the spectrograph HARPS.," The diffuse interstellar band at was analysed using the spectra from GECKO, FEROS, HARPS, CES, UVES and BOES depending on which one was available or had a better signal-to-noise ratio in this range of the spectrum, as shown in One of the set of spectra used to analyse the interstellar absorption features was acquired with the spectrograph HARPS." + High Accuracy Radial velocity Planet Searcher (HARPS) at the ESO La Silla 3.6m telescope is dedicated to the discovery of extra-solar planets and it can measure radial velocities with the highest accuracy currently available., High Accuracy Radial velocity Planet Searcher (HARPS) at the ESO La Silla 3.6m telescope is dedicated to the discovery of extra-solar planets and it can measure radial velocities with the highest accuracy currently available. + HARPS produces the spectra with a resolving power of 115.000 in thespectral range AA.. with a signal-to-noise ratio of 500.," HARPS produces the spectra with a resolving power of 115,000 in thespectral range , with a signal-to-noise ratio of 500." +"where A is a normalization constant, Mpeak is the peak mass of the distribution, and σ is the dispersion.","where A is a normalization constant, $M_{\rm{peak}}$ is the peak mass of the distribution, and $\sigma$ is the dispersion." +" However, since we do not map the peak, the precise parameters of the lognormal function fit are not well constrained, several provide adequate fits to the data points."," However, since we do not map the peak, the precise parameters of the lognormal function fit are not well constrained, several provide adequate fits to the data points." + The function showed in Fig., The function showed in Fig. +" 8 has A-4010, Mpeak=1.55 Mo, and o=0.78."," \ref{mass} has A=4610, $_{peak} =1.55$ $_{\odot}$, and $\sigma=0.78$." +" As argued above, the figure confirms that as a consequence of the nature of the distance distribution, there is relatively little difference in the derived mass distributions whether a single distance is adopted for the clouds or statistical approach is adopted."," As argued above, the figure confirms that as a consequence of the nature of the distance distribution, there is relatively little difference in the derived mass distributions whether a single distance is adopted for the clouds or statistical approach is adopted." +" Also, the results of this analysis are also not strongly dependent on exact parameters of the assumed distance distribution as demonstrated in Appendix A. A number of previous studies have attempted to construct, with samples at least an order of magnitude smaller, the mass distributions of IRDCs (Simon et al."," Also, the results of this analysis are also not strongly dependent on exact parameters of the assumed distance distribution as demonstrated in Appendix A. A number of previous studies have attempted to construct, with samples at least an order of magnitude smaller, the mass distributions of IRDCs (Simon et al." + 2006; Marshall et al., 2006; Marshall et al. + 2009) and fragments within them (Rathborne et al., 2009) and fragments within them (Rathborne et al. + 2006; Ragan et al., 2006; Ragan et al. + 2009)., 2009). + Except for the Ragan et al., Except for the Ragan et al. +" study, the mass distributions in these studies agree: the IRDC mass distribution is similar to that of CO clumps, while the distribution for the sub-structures are steeper, more like the Salpeter IMF."," study, the mass distributions in these studies agree: the IRDC mass distribution is similar to that of CO clumps, while the distribution for the sub-structures are steeper, more like the Salpeter IMF." +" In their analysis of 11 IRDCs, Ragan et al. ("," In their analysis of 11 IRDCs, Ragan et al. (" +"2009) found that the mass distributions of what they calledclumps, which correspond to fragments here, is quite flat, similar to the CO clump mass distribution, in contrast with the present study.","2009) found that the mass distributions of what they called, which correspond to fragments here, is quite flat, similar to the CO clump mass distribution, in contrast with the present study." + However it is difficult to understand the Ragan et al., However it is difficult to understand the Ragan et al. + result as the radii and masses they quote for their clumps imply oopacities over 10 times larger than the oopacities they quote., result as the radii and masses they quote for their clumps imply opacities over 10 times larger than the opacities they quote. +varving cach parameter in order to minimise the X7. value (ος see Jellers&Donati(2008):Marsdenetal. (2005))).,"varying each parameter in order to minimise the $\chi{^2}$ value (eg see \citet{Jeffers08,Marsden05}) )." + These Κον stellar parameters were determined sequentiallv. irstvyaq. thenesinz. and finally the inclination angle of he star.," These key stellar parameters were determined sequentially, first, then, and finally the inclination angle of the star." + ‘This sequence was repeated cach time additional »wameters were modified as a result of the imaging process., This sequence was repeated each time additional parameters were modified as a result of the imaging process. + The full set of parameters that gave the minimum X7 value are shown in Table 3. and were adopted when producing the inal Doppler imaging map in Figure 2((a)., The full set of parameters that gave the minimum $\chi{^2}$ value are shown in Table \ref{parameters} and were adopted when producing the final Doppler imaging map in Figure \ref{Photometry_incl}( (a). + By incorporating he photometry [rom MINO into the imaging process. and were determined. by minimising the. deviations )etween the measured. ancl modelled. data in both the V- and H-bands.," By incorporating the photometry from MKO into the imaging process, and were determined by minimising the deviations between the measured and modelled data in both the V- and R-bands." + The effective photospheric temperature was determined. to be 5900. — 50Ix. which is consistent with hat calculated: based. on the photometric colours of the star anc using the bolometric corrections of Bessellοἱal.(1998) within the error bars of both determinations.," The effective photospheric temperature was determined to be 5900 $\pm$ 50K, which is consistent with that calculated based on the photometric colours of the star and using the bolometric corrections of \citet{Bessell98} within the error bars of both determinations." + The spot temperature was determined to be 4000 + 50Ix. The ohoetospheric-spot. temperature difference of LOOK for this early G-dwarf is consistent with measurements mace by other authors such as Marsdenetal.(2011a.b.2005). ancl supports the relationship developed by Bercyugina(2005).," The spot temperature was determined to be 4000 $\pm$ 50K. The photospheric-spot temperature difference of 1900K for this early G-dwarf is consistent with measurements made by other authors such as \citet{Marsden10a,Marsden10b,Marsden05} and supports the relationship developed by \citet{Berdyugina05}." +. The intensity map with the photometry included is shown in Figure 2((b)., The intensity map with the photometry included is shown in Figure \ref{Photometry_incl}( (b). + There have been a number of dillerent. wavs in which dillerential rotation have been measured., There have been a number of different ways in which differential rotation have been measured. + Reiners&Schmitt(2003) have used a Fourier transform method. to derive the parameters for rapidly. rotating F-tvpe / carly C-type inactive stars., \citet{Reiners03} have used a Fourier transform method to derive the parameters for rapidly rotating F-type / early G-type inactive stars. + They produce a deconvolved line profile from the stellar spectra and then determine the ratio of the second and first zero of the resulting Fourier transform., They produce a deconvolved line profile from the stellar spectra and then determine the ratio of the second and first zero of the resulting Fourier transform. + This ratio is à measure of dillerential rotation on the star., This ratio is a measure of differential rotation on the star. + However. for active voung stars with significant asymmetry within the line profiles. this technique is not as ellective.," However, for active young stars with significant asymmetry within the line profiles, this technique is not as effective." + The differential rotation of a star can also be estimated. by tracking spot features at. cilferent latitudes on the surface of the star., The differential rotation of a star can also be estimated by tracking spot features at different latitudes on the surface of the star. + A solar-like dilferential rotation law. as defined in equation 1.. has been applied in a number of solar-tvpe stars etal.(200328):Alarsclen(2005. 2006))).," A solar-like differential rotation law, as defined in equation \ref{DR}, has been applied in a number of solar-type stars \citet{Donati03a, Marsden05, Marsden06}) )." + where (8) is the rotation rate at latitude 8 in O4 is equatorial rotation rate. and d£? js the rotational shear between the equator and. the pole.," where $\Omega(\theta)$ is the rotation rate at latitude $\theta$ in, $\Omega$ is equatorial rotation rate, and $d\Omega$ is the rotational shear between the equator and the pole." +" In.this paper he dillerential rotation of LID 106506 was measured. using a technique that systematically: adjusts. the differential rotation parameters. Qe, and OQ and determines the best it to the data. by minimising the dillerence between the data and the model fits using X7 minimisation techniques."," Inthis paper the differential rotation of HD 106506 was measured using a technique that systematically adjusts the differential rotation parameters, $\Omega$ and $\Omega$ and determines the best fit to the data, by minimising the difference between the data and the model fits using $\chi^{2}$ minimisation techniques." + This technique has been successfully applied to a number of ate F-f/earlv-G stars (ie. Petit.Donati&CollierCameronal. (2006))).," This technique has been successfully applied to a number of late F-/early-G stars (i.e. \citet{Petit02, Petit04, Barnes05, Marsden06}) )." +" Initially, ao fixed. spot. filling. factor of 0.062. as determined. from the reconstructed. map. was used. and QO. and dQ were considered. to. be free parameters."," Initially, a fixed spot filling factor of 0.062, as determined from the reconstructed map, was used and $\Omega$ and $d\Omega$ were considered to be free parameters." + The magnitude of the dillerential rotation was then determined by fitting a paraboloid to the reduced X7 values., The magnitude of the differential rotation was then determined by fitting a paraboloid to the reduced $\chi^{2}$ values. +" The reduced VO values for the various Q,, and dO is shown in Figure 4..", The reduced $\chi^{2}$ values for the various $\Omega$ and $d\Omega$ is shown in Figure \ref{DR_1}. + The associated error ellipse. as shown in Figure 4. was determined by varving both the spot filling factor and theinclination angle.," The associated error ellipse, as shown in Figure \ref{DR_1} was determined by varying both the spot filling factor and theinclination angle." + Both were changed. by + 10., Both were changed by $\pm$ $\%$. + When varving the inclination angle. the associated multiplications factors were re-determined and the resulting spot. filling [actor as found to match the new anele by minimising the v value.," When varying the inclination angle, the associated multiplications factors were re-determined and the resulting spot filling factor as found to match the new angle by minimising the $\chi^{2}$ value." +" The minimum value for Qe, was determined to be 4.54 + 0.01 land dQ = 0.21οιὃν.", The minimum value for $\Omega$ was determined to be 4.54 $\pm$ 0.01 and $d\Omega$ = $0.21_{-0.03}^{+0.02}$. + Phe errors for the value ofd. the cilferential rotation were determined. by varving both the stars inclination and. spot filling factors bv + and determining Έλι and dQ respectively., The errors for the value of the differential rotation were determined by varying both the star's inclination and spot filling factors by $\pm$ and determining $\Omega$ and $d\Omega$ respectively. + By incorporating the dillerential rotation into the maximunm-entropy image reconstruction. the resulting map is shown in Figure 2.. along with the model fits to the LSD profiles in Figure 3..," By incorporating the differential rotation into the maximum-entropy image reconstruction, the resulting map is shown in Figure \ref{Photometry_incl}, along with the model fits to the LSD profiles in Figure \ref{fit_2}." + Figure 2. shows. on the left. the resulting maximum-entropy image reconstructed map eencrated using the spectroscopic data taken at the AAT.," Figure \ref{Photometry_incl} shows, on the left, the resulting maximum-entropy image reconstructed map generated using the spectroscopic data taken at the AAT." + On the right of this ligure is the reconstructed. image created using both the spectroscopic data and the near-simultaneous photometric data taken at the MIXO., On the right of this figure is the reconstructed image created using both the spectroscopic data and the near-simultaneous photometric data taken at the MKO. + The photometry data has been able to enhance the Iower-LIatitude features., The photometry data has been able to enhance the lower-latitude features. + La cach map. an equatorial period of 1.39 days has been used. which was established during the dillerential rotation analysis and the epoch was set to 2454194.93642361. which was the micelle of the observation run.," In each map, an equatorial period of 1.39 days has been used, which was established during the differential rotation analysis and the epoch was set to 2454194.93642361, which was the middle of the observation run." + Figure 5. shows the V and Ro light. curves coupled =μαith the maximum-entropy. fits that were used to produce the enhanced. DL map., Figure \ref{OC_Graph} shows the V and R light curves coupled with the maximum-entropy fits that were used to produce the enhanced DI map. + Figure 3. shows the mocdeκ. proliles overlaved on the measured LSD profiles., Figure \ref{fit_2} shows the model profiles overlayed on the measured LSD profiles. + Phe model— profiles include the differential rotation parameters., The model profiles include the differential rotation parameters. + Ehe dynamic spectrum ofthe residuals between the observed and =nodelled profiles is shown in Figure 6.., The dynamic spectrum ofthe residuals between the observed and modelled profiles is shown in Figure \ref{Residuals}. + Le appears that the Doppler imagine process. alone with incorporating the differential rotation parameters. has removed the majority of 1ο large scale features with only a small amount of structure μαill apparent in this dynamic spectrum.," It appears that the Doppler imaging process, along with incorporating the differential rotation parameters, has removed the majority of the large scale features with only a small amount of structure still apparent in this dynamic spectrum." + The modelling strategv of Donati&Brown(1907). was used to construct the magnetic field topology on LD 106506., The modelling strategy of \citet{DonatiBrown97} was used to construct the magnetic field topology on HD 106506. + This involved. utilising the spherical harmonic. expansions of the surface magnetic field as implemented by Donatiοἱal. (2006)., This involved utilising the spherical harmonic expansions of the surface magnetic field as implemented by \citet{Donati06}. +. Phe spherical harmonic expansion = 40 was selected as this was the minimum value where any further increase resulted in no further changes in the magnetic opologies., The spherical harmonic expansion = 40 was selected as this was the minimum value where any further increase resulted in no further changes in the magnetic topologies. + Initially. a poloidal plus toroidal Ποια was. assumed. hen the images were created by fitting the data to within he noise level ancl resulted in a mean field. strength. of 68.84 Ci. A similar approach to measuring the cdillerential rotation parameters as was used for the Stokes 7 (brightness) eatures.," Initially, a poloidal plus toroidal field was assumed, then the images were created by fitting the data to within the noise level and resulted in a mean field strength of 68.84 G. A similar approach to measuring the differential rotation parameters as was used for the Stokes $\it{I}$ (brightness) features." + When considering these magnetic features. the equatorial rotation rate Qn was 4.51 + 0.01 +. with a ohotospheric shear 09 of 0.24 + 0.03 1.," When considering these magnetic features, the equatorial rotation rate $\Omega$ was 4.51 $\pm$ 0.01 , with a photospheric shear $\delta\Omega$ of 0.24 $\pm$ 0.03 ." +" ""hisdl implies hat the equatorial rotation period. was d1.38 clavs with a shear approximately 4 times that of the solar value.", This implies that the equatorial rotation period was 1.38 days with a shear approximately 4 times that of the solar value. + TheΤΙ reconstructed.tructecl magneticenetic fieldfields are shown| in liguFigure 7.., The reconstructed magnetic fields are shown in Figure \ref{mapmag}. . + The fits to the Stokes V. LSD profiles are given in Figure, The fits to the Stokes $\it{V}$ LSD profiles are given in Figure \ref{mag_fits}. . +It is generally believed today that we live in an expanding Universe.,It is generally believed today that we live in an expanding Universe. + After the discovery of CMBR (Penzias&Wil-1965) the big-bang cosmology has become the standard model for cosmology which accommodates a beginning of the Universe at some finite past.," After the discovery of CMBR \citep{penin, dicke} the big-bang cosmology has become the standard model for cosmology which accommodates a beginning of the Universe at some finite past." +" However, on its own big-bang cosmology does face some problems both in early and late universe."," However, on its own big-bang cosmology does face some problems both in early and late universe." +" A number of problems cripped up when one describes the early universe, namely, the horizon problem, flatness problem etc."," A number of problems cripped up when one describes the early universe, namely, the horizon problem, flatness problem etc." + The above be resolved evoking a phase of (CurlD1987TORTcSat]OST[d[roinionrecent at observationsa very early epoch.," The above problems can be resolved evoking a phase of inflation \citep{guth,sato,linde,stein} at a very early epoch." + On the [Nbreother hand 1982) predict that our universe is passing through a phase of acceleration , On the other hand recent observations predict that our universe is passing through a phase of acceleration \citep{riess}. +This phase of acceleration is believed to be a late time ∥phase of the universe and may be accommodated in the standard model with a positive cosmological constant., This phase of acceleration is believed to be a late time phase of the universe and may be accommodated in the standard model with a positive cosmological constant. +" Despite its overwhelming success, modern big-bang cosmology still has some unresolved issues."," Despite its overwhelming success, modern big-bang cosmology still has some unresolved issues." +" The physics of the inflation and introduction of a small cosmological constant for late acceleration, is not clearly understood((AIbrecht]2000;[Car]"," The physics of the inflation and introduction of a small cosmological constant for late acceleration, is not clearly understood \citep{reval,sean}." + Thisis why thereis enough motivation to search for froll2001)...alternative cosmology., This is why there is enough motivation to search for alternative cosmology. + Emergent universe (EU) models are employed to find a model which would accommodate the early inflationary phase and avoid the messy situation of the initial singularity ∙∙," Emergent universe (EU) models are employed to find a model which would accommodate the early inflationary phase and avoid the messy situation of the initial singularity \citep{ellis,har}." +" EU scenario can be realized in the framework of general relativity (Mukherjeeetal][2006),, Gauss-Bonnet gravity Brane world gravity (Banerjeej(Paul_&Ghose2010), Brans-Dicke theory (delCampo,Herrera&Labranaj|2007) etc."," EU scenario can be realized in the framework of general relativity \citep{eu}, Gauss-Bonnet gravity \citep{eugb}, Brane world gravity \citep{b1,deb}, Brans-Dicke theory \citep{eubd} etc." + Emergent universe are late time de-Sitter and thus naturally incorporate the late time accelerating phase as well., Emergent universe are late time de-Sitter and thus naturally incorporate the late time accelerating phase as well. + One such model was proposed by (2006) in which a polytropic equation of state (EOS)in the form: where A and B are constants is used., One such model was proposed by \citet{eu} in which a polytropic equation of state (EOS)in the form: where $A$ and $B$ are constants is used. + This is a special case of a more general equation with a=1/2., This is a special case of a more general equation with $\alpha=1/2$. + For such EOS a phenomenological construction can be found in string theories where most of the time models interpolate between two phases of universe (Fabrisetal.||2007)., For such EOS a phenomenological construction can be found in string theories where most of the time models interpolate between two phases of universe \citep{fabris}. +. Universe in this model can stay large enough to avoid quantum gravitational effects even in the very early universe., Universe in this model can stay large enough to avoid quantum gravitational effects even in the very early universe. + Recently studied the viability of this type of model in the light of recent observational data and established bounds on model parameters A and B. It is shown that the best fit valuefor A may be very small but negative although a small positive value is allowed with 95% confidence., Recently \citet{eu1} studied the viability of this type of model in the light of recent observational data and established bounds on model parameters $A$ and $B$ It is shown that the best fit valuefor $A$ may be very small but negative although a small positive value is allowed with $\%$ confidence. + For a viable cosmology the bounds on A and B are determined for some fixed value of K., For a viable cosmology the bounds on $A$ and $B$ are determined for some fixed value of $K$. +" The parameter A, however, appears in the theory as an integration constant and may be to some other value for different initial configuration."," The parameter $K$, however, appears in the theory as an integration constant and may be fixed to some other value for different initial configuration." +fixed[Specific[&Thakur| recently worked with a more model for a small (2011)value of A (Az 0)., \citet{eu2} recently worked with a more specific model for a small value of $A$ $A \approx 0$ ). + In the original work of it was shown that the choice of A drastically changes the matter energy composition of the Universe that eq. , In the original work of \citet{eu} it was shown that the choice of $A$ drastically changes the matter energy composition of the Universe that eq. \ref{eos1}) ) +can mimic., can mimic. + In the present paper we obtain observational (1)bounds on the model parameters B and K for different choices of A as was considered in These choices correspond to very different compositions of cosmic fluid and it would be interesting to see whether realistic cosmologies are permitted for each case since the theory itself puts some constraints over B , In the present paper we obtain observational bounds on the model parameters $B$ and $K$ for different choices of $A$ as was considered in \citep{eu}.. These choices correspond to very different compositions of cosmic fluid and it would be interesting to see whether realistic cosmologies are permitted for each case since the theory itself puts some constraints over $B$ +conditions.,conditions. + Lt is therefore uncertain whether they could be realized as intermediate asvmptotic forms in practice., It is therefore uncertain whether they could be realized as intermediate asymptotic forms in practice. + In a recent. paper. Blandford Begclman (1999) have sought to resolve both these difficulties by claborating on the basic model. adding terms that parametrize the removal of mass. angular momentum and energy by an outflow. which is considered. necessary ο allow accretion to proceed’.," In a recent paper, Blandford Begelman (1999) have sought to resolve both these difficulties by elaborating on the basic model, adding terms that parametrize the removal of mass, angular momentum and energy by an outflow, which is considered necessary `to allow accretion to proceed'." + In their solution the flow is rotating even though ο=5/3. and the Bernoulli parameter can be negative.," In their solution the flow is rotating even though $\gamma=5/3$, and the Bernoulli parameter can be negative." + However. this was again achieved by adjusting the scalings so that essentially no mass reaches the central object.," However, this was again achieved by adjusting the scalings so that essentially no mass reaches the central object." + The purpose of this letter is to investigate the [Luid dynamics of quasi-spherical accretion from an alternative viewpoint which avoids many of the strictures of steady self-similar solutions., The purpose of this letter is to investigate the fluid dynamics of quasi-spherical accretion from an alternative viewpoint which avoids many of the strictures of steady self-similar solutions. + In the theory of thin disces. a consideration of the initial-value problem. in which a cool ancl narrow ring of viscous Iluid orbiting around a central mass spreads racially ancl is accreted. provides an excellent demonstration of the dynamics of disc accretion and is in some wavs more informative than steady solutions.," In the theory of thin discs, a consideration of the initial-value problem, in which a cool and narrow ring of viscous fluid orbiting around a central mass spreads radially and is accreted, provides an excellent demonstration of the dynamics of disc accretion and is in some ways more informative than steady solutions." + La that problem the surface density satisfies a diffusion equation which. depending on the properties of the viscosity. ds. cither linear. in which case the Green function can be obtained analvtically (Lüsst. 1952: Lvnden-Bell Pringle 1974). or non-linear. in which case similarity methods can be used to obtain an exact special solution of the equation (Pringle 1974: Lin Bodenheimer 1982: Lin Pringle LOST: Lvubarskii Shakura LOST).," In that problem the surface density satisfies a diffusion equation which, depending on the properties of the viscosity, is either linear, in which case the Green function can be obtained analytically (Lüsst 1952; Lynden-Bell Pringle 1974), or non-linear, in which case similarity methods can be used to obtain an exact special solution of the equation (Pringle 1974; Lin Bodenheimer 1982; Lin Pringle 1987; Lyubarskii Shakura 1987)." + This self-similar solution accurately ceseribes the behaviour of all solutions of the initial-value problem long alter the inner radius of the Uuiel first reaches the central object., This self-similar solution accurately describes the behaviour of all solutions of the initial-value problem long after the inner radius of the fluid first reaches the central object. + In. the case of quasi-spherical accretion. the governing equations are non-linear ancl similarity methods can again be used to provide an exact solution which may similarly be expected to be an attractor for the initial-value problem.," In the case of quasi-spherical accretion, the governing equations are non-linear and similarity methods can again be used to provide an exact solution which may similarly be expected to be an attractor for the initial-value problem." + A simplified. one-dimensional model. for non-relativistic quasi-spherical accretion ds. adopted. in) which physical quantities depend only on the spherical radius 7 and on time ἐν, A simplified one-dimensional model for non-relativistic quasi-spherical accretion is adopted in which physical quantities depend only on the spherical radius $r$ and on time $t$. + Phe model is described by the equations and Llere p is the densitv. 4 the radial velocity. © the angular velocity. |?= GAL/r the gravitational potential (where AJ is the mass of the central object). p the pressure. f—a7O the specific angular momentum. fj the (dynamic) viscosity. (which. for simplicity. is assumed. to act only on the differential rotation) and e the specific internal energy.," The model is described by the equations and Here $\rho$ is the density, $u$ the radial velocity, $\Omega$ the angular velocity, $\Phi=-GM/r$ the gravitational potential (where $M$ is the mass of the central object), $p$ the pressure, $\ell=r^2\Omega$ the specific angular momentum, $\mu$ the (dynamic) viscosity (which, for simplicity, is assumed to act only on the differential rotation) and $e$ the specific internal energy." + The Lagrangian time derivative is lt is further assumed that the [uid is an ideal gas with and that the viscosity is given. as in NY. by an a-mocel. where Og=(Αν1Ξ is the Weplerian angular velocity.," The Lagrangian time derivative is It is further assumed that the fluid is an ideal gas with and that the viscosity is given, as in NY, by an $\alpha$ -model, where $\Omega_{\rm K}=(GM/r^3)^{1/2}$ is the Keplerian angular velocity." + “Phe parameters 5? anc à are assumed to be constant., The parameters $\gamma$ and $\alpha$ are assumed to be constant. + These equations may be obtained from the full axisvmumetric equations in spherical polar coordinates (0.0.0) bv considering. the equatorial plano @=x/2 and neglecting terms associated with the 8-velocitv and any 8-dependence.," These equations may be obtained from the full axisymmetric equations in spherical polar coordinates $(r,\theta,\phi)$ by considering the equatorial plane $\theta=\pi/2$ and neglecting terms associated with the $\theta$ -velocity and any $\theta$ -dependence." + “Lhe variables may be considered as approximate spherically averagecl quantities., The variables may be considered as approximate spherically averaged quantities. + Although approximate. the equations have two important. properties: they reduce exactly to those of spherical aceretion in the absence of rotation. and they possess exact conservation laws for mass. angulare momentum and energy.e which are of the form and where is the specific total energy ancl the total energy Lux.," Although approximate, the equations have two important properties: they reduce exactly to those of spherical accretion in the absence of rotation, and they possess exact conservation laws for mass, angular momentum and energy, which are of the form and where is the specific total energy and the total energy flux." + Phese equations are appropriate to quasi-spherical accretion. rather than disc accretion. because the dissipated energy is retained in the [uid and may be advected with the Dow.," These equations are appropriate to quasi-spherical accretion, rather than disc accretion, because the dissipated energy is retained in the fluid and may be advected with the flow." + NY solved. essentially equations (1)) (4)) for. the case of a steady. racially self-similar Wow.," NY solved essentially equations \ref{drho}) \ref{de}) ) for the case of a steady, radially self-similar flow." + Such a solution is unbouncled in size. mass. angular momentum ancl energy. and must therefore be interpreted. as an approximation to a finite steady solution within a region [ar removed [rom its boundaries.," Such a solution is unbounded in size, mass, angular momentum and energy, and must therefore be interpreted as an approximation to a finite steady solution within a region far removed from its boundaries." + In contrast. let us seek a time-dependent," In contrast, let us seek a time-dependent" +2nd observation.,2nd observation. + This is not the case for the 3rd observation; errors of photon index and absorption column density are large., This is not the case for the 3rd observation; errors of photon index and absorption column density are large. +" When we fixed the photon index to 2.0 for the 3rd observation, the absorption column density of the uniform absorber becomes reasonable around 1x1073 cm~?."," When we fixed the photon index to 2.0 for the 3rd observation, the absorption column density of the uniform absorber becomes reasonable around $1\times10^{23}$ $^{-2}$." + The fitting results are summarized in table 6.., The fitting results are summarized in table \ref{fitdiff}. +" The partial absorber has a large column density of >1024 cm""? with a large covering fraction of >0.3, and they seem variable."," The partial absorber has a large column density of $>10^{24}$ $^{-2}$ with a large covering fraction of $>0.3$, and they seem variable." + These behaviors could cause a complex time variability., These behaviors could cause a complex time variability. + Such a Compton-thick absorber was for the first time observed for the Cen A. The partial covering absorber is thought to be variable during each observation., Such a Compton-thick absorber was for the first time observed for the Cen A. The partial covering absorber is thought to be variable during each observation. +" Therefore, next, we investigated a spectral variability with a finer time resolution."," Therefore, next, we investigated a spectral variability with a finer time resolution." +" We divided one observation into several periods with a step of 10 ks, following the light curve in figure 2,, and took a difference spectrum between two periods, one of which is the low flux period just before brightening and the other is around the highest flux level."," We divided one observation into several periods with a step of 10 ks, following the light curve in figure \ref{lc}, and took a difference spectrum between two periods, one of which is the low flux period just before brightening and the other is around the highest flux level." +" We took the difference between 7th and 3rd bin during the 1st observation, 10th and 9th during the 1st observation, 24th and 16th bin during the 2nd observation, 41th and 38th bin during the"," We took the difference between 7th and 3rd bin during the 1st observation, 10th and 9th during the 1st observation, 24th and 16th bin during the 2nd observation, 41th and 38th bin during the" +since σι is within (he near zone. we can use (he expansions where. again. (0) means vu= 0.,"Since $\sigma_1$ is within the near zone, we can use the expansions where, again, $(0)$ means $u=0$ ." +" Recalling that k-d,(0)=0. and substituting these expressions into Eq. (28))."," Recalling that ${\bf k} \cdot {\bf +d}_a(0)=0$, and substituting these expressions into Eq. \ref{V1}) )," + we obtain Combining-∙− the expressions⋅ ∙ (24))↜↜ ancl . (30))MN [or. VC9] and Wi.1) we see that all terms involving σι cancel. as expected.," we obtain Combining the expressions \ref{V2final}) ) and \ref{V1final}) ) for $\Psi_a^{(2)}$ and $\Psi_a^{(1)} $, we see that all terms involving $\sigma_1$ cancel, as expected." +" Our final expression for V, is Notice that the dependence on e, in the term K-v,(0)/c, in Eq. (31))", Our final expression for $\Psi_a$ is Notice that the dependence on $c_g$ in the term ${\bf k} \cdot {\bf v}_a(0) /c_g$ in Eq. \ref{Vfinal}) ) +" is illusory: to obtain the measured effect. we must multiply V, for body à by its mass min. and sum over all the bodies in the svstem."," is illusory: to obtain the measured effect, we must multiply $\Psi_a$ for body $a$ by its mass $m_a$, and sum over all the bodies in the system." +" Although the event corresponding to &=0 (closest approach) is different for each body. the effect of these differences on v, is of order of the acceleration ol each body. which is O(e). hence we can evaluate v,(0) at Che same coordinate (time lor each body. sav. the time of reception of the ray."," Although the event corresponding to $\sigma=0$ (closest approach) is different for each body, the effect of these differences on $v_a$ is of order of the acceleration of each body, which is $O(\epsilon)$, hence we can evaluate ${\bf v}_a(0)$ at the same coordinate time for each body, say, the time of reception of the ray." +" But. because $7,m,v,=0 in barveentric coordinates. the net elfect of this c;-dependent term vanishes."," But, because $\sum_a m_a {\bf +v}_a = 0$ in barycentric coordinates, the net effect of this $c_g$ -dependent term vanishes." +" Sinilaulv. for a given light rav. (he lime a, is a constant. independent of the body."," Similarly, for a given light ray, the time $\sigma_e$ is a constant, independent of the body." +" The difference between (he various values of s. lor different bodies has an effect only at order e. so again. summing over all the bodies in the system: causes the velocity correction to the In|o,| term to vanish."," The difference between the various values of $s_e$ for different bodies has an effect only at order $\epsilon$, so again, summing over all the bodies in the system causes the velocity correction to the $\ln |\sigma_e|$ term to vanish." +" The Inσσ, term (hen is an irrelevant constant.", The $\ln |c\sigma_e|$ term then is an irrelevant constant. + By the same argument.the factor of 2 in the logarithm produces only an irrelevant constant.," By the same argument,the factor of 2 in the logarithm produces only an irrelevant constant." + The final formula for the Gime delay. correct to 1208 order is," The final formula for the time delay, correct to 1.5PN order is" +and the lack of any proximity effect near the QSO.,and the lack of any proximity effect near the QSO. +" If ionizing radiation were escaping from the AGN, the surrounding IGM should be ionized within 5-10 Mpc of the QSO."," If ionizing radiation were escaping from the AGN, the surrounding IGM should be ionized within 5–10 Mpc of the QSO." +" The aabsorption extends Az£z0.019 beyond the previous QSO redshift estimate, zem=2.885 (Reimers "," The absorption extends $\Delta z \approx 0.019$ beyond the previous QSO redshift estimate, $z_{\rm em} = 2.885$ (Reimers 1997)." +"As discussed in Section 3.1, this dilemma is easily resolved11997). by moving the quasar redshift to z=2.904, an offset (1460 aat z= 2.9) that corresponds to 19.5 Mpc comoving proper distance."," As discussed in Section 3.1, this dilemma is easily resolved by moving the quasar redshift to $z = 2.904$, an offset (1460 at $z = 2.9$ ) that corresponds to 19.5 Mpc comoving proper distance." +" However, the two proposed systemic redshifts, zag0=2.885 and 2.904, raise several questions."," However, the two proposed systemic redshifts, $z_{\rm QSO} = 2.885$ and 2.904, raise several questions." + Does the QSO lie inside an aand ccavity produced by its photoionizing radiation?, Does the QSO lie inside an and cavity produced by its photoionizing radiation? +" Or, is the QSO ionizing radiation absorbed by circumnuclear gas?"," Or, is the QSO ionizing radiation absorbed by circumnuclear gas?" + Is there infalling gas near the AGN that produces associated absorption with zaps>Zem?, Is there infalling gas near the AGN that produces associated absorption with $z_{\rm abs} > z_{\rm em}$? + It is worth investigating these possibilities., It is worth investigating these possibilities. +" For ((1 ryd) and ((4-ryd) ionizing-photon luminosities, si) and ο, what Strómmgren radii will this bright QSO carve out?"," For (1 ryd) and (4-ryd) ionizing-photon luminosities, $S_0^{\rm (H)}$ and $S_0^{\rm (He)}$, what Strömmgren radii will this bright QSO carve out?" +" These distances are obviously limiting cases, in ionization equilibrium, whereas one expects the ionization zones around the QSO to involve time-dependent propagation of ionization fronts."," These distances are obviously limiting cases, in ionization equilibrium, whereas one expects the ionization zones around the QSO to involve time-dependent propagation of ionization fronts." +" Indeed, the lack of a proximity effect could arise from this source just starting to burn its way out of the galactic nucleus."," Indeed, the lack of a proximity effect could arise from this source just starting to burn its way out of the galactic nucleus." +" Given the strength of the z=2.9 absorber, we will assume higher gas"," Given the strength of the $z = 2.9$ absorber, we will assume higher gas" +size and a lower limit for the magnification lactor that our similar to ours.,size and a lower limit for the magnification factor that our similar to ours. + Accounting for the influence οἱ &ravitational. lensing. the velocity-integrated CO(1-0) Dux density implies that the nuclear content of molecular gas in is ~10hFALL. assuming ao CO-to-Lls conversion factor of ~LAL.kmsec.+per) which is tvpical for starburstkinematically violent. UV-intense. environments of gas-rich. Ht-ultraluminous svstems (Downes Solomon 1998).," Accounting for the influence of gravitational lensing, the velocity-integrated CO(1-0) flux density implies that the nuclear content of molecular gas in is $\sim10^{10} h^{-2} M_\odot$, assuming a $_2$ conversion factor of ${\rm\sim 1 (M_\odot\ km\ +sec^{-1}\ pc^2)^{-1}}$ which is typical for starburst/kinematically violent, UV-intense environments of gas-rich, IR-ultraluminous systems (Downes Solomon 1998)." +" Assuming the CO is in a rotating disk. the dynamical mass can be calculated from the radius of 2500 pc set bv the lens modeling. anc using a rotational velocity of 350 km | set by the observed. of line velocity LINLIM = 250 km n and assuming. a disk. inclination.. angle of 45"" (Downes et al."," Assuming the CO is in a rotating disk, the dynamical mass can be calculated from the radius of $\sim500$ pc set by the lens modeling, and using a rotational velocity of 350 km $^{-1}$ set by the observed of line velocity HWHM = 250 km $^{-1}$ and assuming a disk inclination angle of $^o$ (Downes et al." + 1999)., 1999). + Ehe implied. dynamical mass is 1.5101 NM. within ~500 pe of the nucleus. consistent with the value derived from the CO Dux.," The implied dynamical mass is $1.5\times 10^{10}$ $_\odot$ within $\sim$ 500 pc of the nucleus, consistent with the value derived from the CO flux." + Hence it appears likely that the molecular gas mass makes a significant. and perhaps domünant. contribution to the total mass within a [ew huncdred: parsees of the nucleus in5255... unless the nuclear €O disk is close to face-on.," Hence it appears likely that the molecular gas mass makes a significant, and perhaps dominant, contribution to the total mass within a few hundred parsecs of the nucleus in, unless the nuclear CO disk is close to face-on." + A similar conclusion has been reached for most nearby nuclear starburst galaxies (Downes ancl Solomon 1998)., A similar conclusion has been reached for most nearby nuclear starburst galaxies (Downes and Solomon 1998). + This paper has presented resolved: images of nuclear CO(1-0) emission in the gravitationally lensecl BAL quasar5255., This paper has presented resolved images of nuclear CO(1-0) emission in the gravitationally lensed BAL quasar. +.. While the continuum emission is found o be well aligned with the optical quasar images. the CO(I-) is more extended. with a broken ring-like appearance.," While the continuum emission is found to be well aligned with the optical quasar images, the CO(1-0) is more extended, with a broken ring-like appearance." + Such a structure is consistent with the action of gravitational ensing. with the continuum emission occurring on the scale of the quasar core. while the CO(1-0). arises from à arger region and is cilferentiallv magnilied.," Such a structure is consistent with the action of gravitational lensing, with the continuum emission occurring on the scale of the quasar core, while the CO(1-0) arises from a larger region and is differentially magnified." + Phe threc-image nature of thas posed a problem for lens modeling. as an extremely arge. [lat core is required. to. produce the central image.," The three-image nature of has posed a problem for lens modeling, as an extremely large, flat core is required to produce the central image." + Such three image configurations are a nature consequence of gravitational lensing by a fattened potential which can ooduce naked. cusps., Such three image configurations are a nature consequence of gravitational lensing by a flattened potential which can produce naked cusps. + Modeling of the ςΟ(1-0) emission supports this hypothesis. although a deficit in constraints implies that the model is not unique.," Modeling of the CO(1-0) emission supports this hypothesis, although a deficit in constraints implies that the model is not unique." + An immecdiate »edietion of this model is that the lensing galaxy. whose »oxition. could. be revealed by observing below the Lyman imit for this system {4400). hence removing the elare rom the quasars. should: be offset. .5aresec. [rom the quasar image. rather than lying behind the quasar images.," An immediate prediction of this model is that the lensing galaxy, whose position could be revealed by observing below the Lyman limit for this system ${\rm (\lta 4400\AA)}$, hence removing the glare from the quasars, should be offset $\sim0.5$ arcsec from the quasar image, rather than lying behind the quasar images." + Currently. our CO images of are. of limited signal-to-noise.," Currently, our CO images of are of limited signal-to-noise." + Llowever. with further integration a detailed map of the CO image can be mace.," However, with further integration a detailed map of the CO image can be made." + As this region will be free from the elfects of microlensing. and as its extended nature provides many more constraints (Ixochanek. Keeton MeLoed 2001). such imaging has the potential to provide a more accurate model of the lensing in tthan from the quasar images.," As this region will be free from the effects of microlensing, and as its extended nature provides many more constraints (Kochanek, Keeton McLoed 2001), such imaging has the potential to provide a more accurate model of the lensing in than from the quasar images." +"where nj= defe. ca=OFfe, is the Alfvenn speed. zi; is the ton plasma frequency of the background plasma.","where $n_\parallel=k_\parallel c/\omega$ , $v_A=\Omega_i/\omega_{pi}$ is the Alfvénn speed, $\omega_{pi}$ is the ion plasma frequency of the background plasma." + Using the neutralization conditions (4)). (6)) can be rewritten in the form The Cl. components are where with (rove)—blui. lus). 6=Oshfe and {4Gr) the step function.," Using the neutralization conditions \ref{eq:neutral}) ), \ref{eq:K12a}) ) can be rewritten in the form The CR components are where with $(x_1,x_2)=b(1/u_1,1/u_2)$ , $b=\Omega/|k_\parallel|c$ and $H(x)$ the step function." + Eq (9)) and (10)) correspond. to the clissipative and reactive part respectively of the response tensor due to Clits., Eq \ref{eq:chia}) ) and \ref{eq:chih}) ) correspond to the dissipative and reactive part respectively of the response tensor due to CRs. + In the strong magnetic field limit cz1. the clissipative part is zero.," In the strong magnetic field limit $x>1$, the dissipative part is zero." + When we«1. both dissipative and reactive terms contribute to the dispersion.," When $x<1$, both dissipative and reactive terms contribute to the dispersion." + The dissipative part is due to CRs in the resonance e—ησεµ as discussed extensively in the context of the resonant instability (see. cliscussion below) (κακιά&Pearce1969:MelroseWentzel1970: 19583a).," The dissipative part is due to CRs in the resonance $x\equiv b/u\approx \mu$ as discussed extensively in the context of the resonant instability (see discussion below) \citep{kp69,mw70,s75,mv82,lc83a}." +". The dispersion relation takes the following form Inthe Χο= Oanedy, »0 limit (LU must also be zero). (12)) with the lower sign reproduces the usual Alfvénn moce dispersion w=|é||e and with the upper sign the [ast mode dispersion w=key."," The dispersion relation takes the following form In the $\chi_a\to0$ and $\chi_h\to 0$ limit $\chi^{(0)}_h$ must also be zero), \ref{eq:disp}) ) with the lower sign reproduces the usual Alfvénn mode dispersion $\omega=|k_\parallel| v_A$ and with the upper sign the fast mode dispersion $\omega=kv_A$." + Instability. occurs. when a has an imaginary part of the appropriate sign., Instability occurs when $\omega$ has an imaginary part of the appropriate sign. + An imaginary part can arise in two cillerent wavs. which we refer to as resonant and nonresonant instabilities.," An imaginary part can arise in two different ways, which we refer to as resonant and nonresonant instabilities." + Writing w-w&|UE in (12)) and assuming EL«uw. the resonant instability is described. by equating the small imaginary terms. giving --IRE.," Writing $\omega\to\omega+i\Gamma$ in \ref{eq:disp}) ) and assuming $\Gamma\ll\omega$, the resonant instability is described by equating the small imaginary terms, giving $\Gamma=\chi_ak_\parallel^2v_A^2/2\omega$." +" The value of x, is determined by the Ch distribution.", The value of $\chi_a$ is determined by the CR distribution. + A resonant instability can develop for both Alfvénn ancl fast-mode waves if streaming CRs are in resonance with the wave considered., A resonant instability can develop for both Alfvénn and fast-mode waves if streaming CRs are in resonance with the wave considered. + For Alfvénn waves (v=LE ry). the growth rate 15 The polarization of both modes is generally elliptical ancl becomes approximately circular. for nearly parallel propagation [A1/&|«1.," For Alfvénn waves $\omega=|k_\parallel|v_A$ ), the growth rate is The polarization of both modes is generally elliptical and becomes approximately circular for nearly parallel propagation $|k_\perp/k_\parallel|\ll1$." + In each modo. waves of both sense of polarization can grow.," In each mode, waves of both sense of polarization can grow." + On neglecting χα. (12)) with the upper sign becomes a real equation. because the quantity inside the square root is a sum of squares.," On neglecting $\chi_a$, \ref{eq:disp}) ) with the upper sign becomes a real equation, because the quantity inside the square root is a sum of squares." + For the negative sign (AMfvénn waves). u can be negative. ancl one of the two solutions corresponds to an intrinsically erowing wave.," For the negative sign (Alfvénn waves), $\omega^2$ can be negative, and one of the two solutions corresponds to an intrinsically growing wave." + This is identified as the nonresonant instability., This is identified as the nonresonant instability. + For |kifA|«1. the condition for this nonresonant instabilitv to occur is LUBUyo1.," For $|k_\perp/k_\parallel|\ll1$, the condition for this nonresonant instability to occur is $(\chi^{(0)}_h+\chi_h)^2>1$." + Phe erowth rate is In principle. this can be satisfied either due to |i]71 or to [yy|oL.," The growth rate is In principle, this can be satisfied either due to $|\chi_h|>1$ or to $|\chi_h^{(0)}|>1$ ." + For &=1. the integral v; given by (10)) can be calculated exactly. Le. The second equality in (15)) is obtained for p=2.," For $\sigma=1$, the integral $\chi_h$ given by \ref{eq:chih}) ) can be calculated exactly, i.e., The second equality in \ref{eq:chih2}) ) is obtained for $p=2$." + The condition [v|] cannot be satisfied for wryl., The condition $|\chi_h|>1$ cannot be satisfied for $x_1\ll1$. + Thus. only the second. possibilityepe (ALTIED> 1) is. relevant.," Thus, only the second possibility $\chi^{(0)}_h>1$ ) is relevant." + For re-1. one finds so that the contribution from the CR current is much smaller than the contribution due to the compensating current in the background. plasma.," For $x_2\ll x_1\ll 1$, one finds so that the contribution from the CR current is much smaller than the contribution due to the compensating current in the background plasma." + Lt follows that in treating the nonresonant instabilitv. one can neglect. the direct. contribution of the CRs. described by xí. in comparison with the indirect. contribution. described. by iOTEE," It follows that in treating the nonresonant instability, one can neglect the direct contribution of the CRs, described by $\chi_h$, in comparison with the indirect contribution, described by $\chi_h^{(0)}$." +" The growth rate (14)) can be written in the following approximate form Nonresonant instability requires An explanation for why the direct contribution from the CR current is so much smaller than the indirect contribution rom the background plasma is that in the limit Ayr,71. most ofCRs move rather rigicly."," The growth rate \ref{eq:Gam}) ) can be written in the following approximate form Nonresonant instability requires An explanation for why the direct contribution from the CR current is so much smaller than the indirect contribution from the background plasma is that in the limit $k_\parallel r_g\gg1$, most of CRs move rather rigidly." + Pheir only role in this limit is to induce the compensating current in the background ralasma that drives the instability., Their only role in this limit is to induce the compensating current in the background plasma that drives the instability. + Itis appropriate to point out here that a small fraction ofCRs can satisfy the resonant condition (j£— c) and that resonant instability due to these resonant CRs is insignificant compared tothe nonresonant instability due to the compensatingcurrent., Itis appropriate to point out here that a small fraction of CRs can satisfy the resonant condition $\mu=x$ ) and that resonant instability due to these resonant CRs is insignificant compared tothe nonresonant instability due to the compensatingcurrent. + However. the resonant interactions between these Clts and waves with," However, the resonant interactions between these CRs and waves with" +What future work is necded to better understand the ircjonization?,What future work is needed to better understand the reionization? + First. we need to assess whether the HE 1312 sieht line is typical.," First, we need to assess whether the HE $-$ 4342 sight line is typical." + The COS-CTO eai will be observing two more AAQN targets. HSITOO|6116 and Q0302-003. which nay clarity whether the patchy reionization at +«2.8 is cohbunon to most seht lines.," The COS-GTO team will be observing two more AGN targets, HS1700+6416 and Q0302-003, which may clarify whether the patchy reionization at $z < 2.8$ is common to most sight lines." + Second. along the current sieht line. we have ideutified many regious of partial flux rausnussion aud three loug troughs of absorption. where surveys of ealaxies and ACN wil welp clavify the sources of ionization.," Second, along the current sight line, we have identified many regions of partial flux transmission and three long troughs of absorption, where surveys of galaxies and AGN will help clarify the sources of ionization." + Third. iufrare spectra of the UL) A5007. 1959 narrow ciuission ines (at 1.955 yon and 1.936 pau) could confirm the QSO systemic redslift.," Third, infrared spectra of the ] $\lambda$ 5007, 4959 narrow emission lines (at 1.955 $\mu$ m and 1.936 $\mu$ m) could confirm the QSO systemic redshift." + Piuuiug down is eruca or undoerstaudius the nature of the ionizationtose state of he eas near the QSO. which shows no evidence of a xoxinuitv effect.," Pinning down $z_{\rm QSO}$ is crucial for understanding the nature of the ionization state of the gas near the QSO, which shows no evidence of a proximity effect." + Ou the theoretical frout. cosmologica simulations with time-dependent IT aud We chemistiv and Lfrout radiative transfer should clucidate the structure topology of the laments aud voids that we observe iu the Lya--forest aud ttroughs.," On the theoretical front, cosmological simulations with time-dependent H and He chemistry and I-front radiative transfer should elucidate the structure topology of the filaments and voids that we observe in the -forest and troughs." + Serious studies of rrelonization have been lincdered by the lack of maux ACN targets sufficicutly bright to obtain faa-UV spectra., Serious studies of reionization have been hindered by the lack of many AGN targets sufficiently bright to obtain far-UV spectra. + The spectrograpls on aand hhave made good progress on three ACN sight lines. but newly discovered GOP targets (Syplers 20092.b) are considerably fainter than WE 2317. 12312 or | 6116.," The spectrographs on and have made good progress on three AGN sight lines, but newly discovered GP targets (Syphers 2009a,b) are considerably fainter than HE $-$ 4342 or $+$ 6416." + New observations of these falter targets will probably measure only broad-baud ooptical depths. for which the high-resolution COS spectra nav provide euidance i interpretation.," New observations of these fainter targets will probably measure only broad-band optical depths, for which the high-resolution COS spectra may provide guidance in interpretation." + The long-term study of the rreionization process. from ;=2.9.3.3. must await the construction of a future larec-aperture. far-UV telescope in space (Shull 1999: Postiman 22009).," The long-term study of the reionization process, from $z = 2.3-3.3$, must await the construction of a future large-aperture, far-UV telescope in space (Shull 1999; Postman 2009)." + It would be especially useful if this mission had high scusitivity aud moderate spectral resolution 20.000) extending from 1130 ddown to the 912 lydrogen Lyman Lint.," It would be especially useful if this mission had high sensitivity and moderate spectral resolution $R \ga 20,000$ ) extending from 1130 down to the 912 hydrogen Lyman Limit." + This would euable the study. of patchy weionization iu A301 down to redshifts 22., This would enable the study of patchy reionization in $\lambda 304$ down to redshifts $z \ga 2$. + It is our pleasure το ackuowledge the thousauds of people who imade IST Servicing Mission Ll a hnuge success., It is our pleasure to acknowledge the thousands of people who made HST Servicing Mission 4 a huge success. + We thank Brian Necney. Steve Penton. Stépphane Bélland. and the rest of the COS/CTO team for their work on the calibration and verification of the carly COS data. and Cristina Oliviera for her helpful input on CLLlOL seement-B CALCOS processing.," We thank Brian Keeney, Steve Penton, Stépphane Bélland, and the rest of the COS/GTO team for their work on the calibration and verification of the early COS data, and Cristina Oliviera for her helpful input on G140L segment-B CALCOS processing." + We thank Dieter Reimers for discussions regarding the QSO systemic redshift and for providing VLT data on the fforest., We thank Dieter Reimers for discussions regarding the QSO systemic redshift and for providing VLT data on the forest. + Mark Cüroux. Blair Savage. David Syphers. aud the referee provided numerous conuueuts that inproved the arguinenuts made in this paper.," Mark Giroux, Blair Savage, David Syphers, and the referee provided numerous comments that improved the arguments made in this paper." +" This work is based ou observations made with the NASA/ESATelescope, obtained from the data archive at the Space Telescope Science. Iustitute."," This work is based on observations made with the NASA/ESA, obtained from the data archive at the Space Telescope Science Institute." +. STScI is operated by the Association of Universities for Researcli in Astrouoniv. Tuc. uuder NASA contract NÀS5-26555.," STScI is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS5-26555." + Our work was also supported by NASA erants NNNOSACLIG and NAÀS5-O8013. and the Astroplivsical Theory Program (NNNOT-ACT7C from NASA aud ASTO7-07lL from NSF) at the University of Colorado at Boulder., Our work was also supported by NASA grants NNX08AC146 and NAS5-98043 and the Astrophysical Theory Program (NNX07-AG77G from NASA and AST07-07474 from NSF) at the University of Colorado at Boulder. + Our work was also supported by NASA erants NNNOSACLIG and NAÀS5-O8013. and the Astroplivsical Theory Program (NNNOT-ACT7C from NASA aud ASTO7-07lL from NSF) at the University of Colorado at Boulder.L, Our work was also supported by NASA grants NNX08AC146 and NAS5-98043 and the Astrophysical Theory Program (NNX07-AG77G from NASA and AST07-07474 from NSF) at the University of Colorado at Boulder. + Our work was also supported by NASA erants NNNOSACLIG and NAÀS5-O8013. and the Astroplivsical Theory Program (NNNOT-ACT7C from NASA aud ASTO7-07lL from NSF) at the University of Colorado at Boulder.L7, Our work was also supported by NASA grants NNX08AC146 and NAS5-98043 and the Astrophysical Theory Program (NNX07-AG77G from NASA and AST07-07474 from NSF) at the University of Colorado at Boulder. +transverse. and normal componcuts: F=RE|T0Nz.,"transverse, and normal components: $\mathbf{F} = R \mathbf{\hat{r}} + T \mathbf{\hat{\theta}} + N \mathbf{\hat{z}}$." +" The rate of change of seunianajor axis is: where fois the true anomaly, ο is eccentricity. aud AL, aud are the stellar auc planetary niasses."," The rate of change of semi-major axis is: where $f$ is the true anomaly, $e$ is eccentricity, and $M_\star$ and $m_p$ are the stellar and planetary masses." +" The thrust is im,produced by the dipole of the reradiated fiux distribution. which we characterize by dts fractional wuplitude fy. so that the fux distribution has au angular dependence x(L1ficose). where o ds measured from a pole through the peak of the dipole."," The thrust is produced by the dipole of the reradiated flux distribution, which we characterize by its fractional amplitude $f_d$, so that the flux distribution has an angular dependence $\propto (1+f_d \cos \psi)$, where $\psi$ is measured from a pole through the peak of the dipole." + The incomuug starlieht is Doppler-boosted in the frame of the planet. and upon absorption causes angular momentum loss (Povutine-Robertson drag) but that effect is efe—510 tweaker. where e is the orbital velocity and e is the speed of light.," The incoming starlight is Doppler-boosted in the frame of the planet, and upon absorption causes angular momentum loss (Poynting-Robertson drag), but that effect is $v/c \sim 5\times10^{-4}$ weaker, where $v$ is the orbital velocity and $c$ is the speed of light." +" The force due to the dipole courponeut Is where £, is the stars Iundnosity. A, is the plauct’s radius as mieasured by photometry during transit. r=atlολα|ecosf) Lis the distance between tlie planet and star. aud A dis the planets albedo. ("," The force due to the dipole component is where $\mathcal{L}_\star$ is the star's luminosity, $R_p$ is the planet's radius as measured by photometry during transit, $r = a(1-e^2)(1+e \cos f)^{-1}$ is the distance between the planet and star, and $A$ is the planet's albedo. (" +We assuue that the albedo does not vary over the surface ofthe planet: if it does. there is additional opportunity for ποιοται transfer between starlight and the orbit.),"We assume that the albedo does not vary over the surface of the planet; if it does, there is additional opportunity for momentum transfer between starlight and the orbit.)" + We assume atmospheric dvuaimics displace the dipole a longitudinal anele Ay cast of the substellar point. so the compoucuts of this force are: R=FeosAy. T=FnAy. aud NV= 0.," We assume atmospheric dynamics displace the dipole a longitudinal angle $\lambda_d$ east of the substellar point, so the components of this force are: $R = F \cos \lambda_d$, $T = F \sin \lambda_d$, and $N = 0$ ." + For an order-ofinaenitude estimate. let us assiunie Ay aud fa ave coustaut.," For an order-of-magnitude estimate, let us assume $\lambda_d$ and $f_d$ are constant." + There may be an auti-correlatiou iu the magnitudes of Ay and f; as hieh-altitude opacity sources may cause planets with high effective temperature to reracliate the euergv before it is advected (?).., There may be an anti-correlation in the magnitudes of $\lambda_d$ and $f_d$ as high-altitude opacity sources may cause planets with high effective temperature to reradiate the energy before it is advected \citep{2007Fb}. + Averaging equation (1)) over time. we have: Therefore the timescale for seii-miajor axis change is: where reference values are for Jupiter aud the Sun.," Averaging equation \ref{eq:dadt_md}) ) over time, we have: Therefore the timescale for semi-major axis change is: where reference values are for Jupiter and the Sun." + Therefore the timescale of orbital chauge for a typical hot Jupiter is about 100 times its age., Therefore the timescale of orbital change for a typical hot Jupiter is about $100$ times its age. +" The rate of change in ecceutrityv is: where € is the eccentric anomaly,", The rate of change in eccentrity is: where $\mathcal{E}$ is the eccentric anomaly. + Now. using tle samc force and asstuptions as above. we have: where the final approximation ids first-order in eccentricity.," Now, using the same force and assumptions as above, we have: where the final approximation is first-order in eccentricity." + The Yarkovsky effect causes eccentricity to erow if the hotspot is shifted to the east (c.g... 2)). on a timescale simular to that of scuai-major axis erowth.," The Yarkovsky effect causes eccentricity to grow if the hotspot is shifted to the east (e.g., \citealt{2002SG}) ), on a timescale similar to that of semi-major axis growth." + The assuniptious of constant fy and A; will break down at laree eccentricity. which will be shown iu refsec:calc..," The assumptions of constant $f_d$ and $\lambda_d$ will break down at large eccentricity, which will be shown in \\ref{sec:calc}." + However. given that the Yarkovsky effect will ouly be large for siiall values of Πα. tidal dissipation in the planet should be sufficicut to damp the eccentricity faster than the Yarkovskw effect excites it.," However, given that the Yarkovsky effect will only be large for small values of $R_p/a$, tidal dissipation in the planet should be sufficient to damp the eccentricity faster than the Yarkovsky effect excites it." + The rate of change in orbital inclination is: £xΑν , The rate of change in orbital inclination is: $\dot I \propto N$. +Above we asstuued that the reradiation field is anisotropic. but is still sauinetrie above aud below the plane of the orbit.," Above we assumed that the reradiation field is anisotropic, but is still symmetric above and below the plane of the orbit." + Iu reality. the normal component of the force. NV. fluctuates with fluctuating eddies.," In reality, the normal component of the force, $N$, fluctuates with fluctuating eddies." + Towever. this is a raucdom walk with extremely small step size. so it is unlikely to ect very far. even if the normalization of the rate were large.," However, this is a random walk with extremely small step size, so it is unlikely to get very far, even if the normalization of the rate were large." + Iu this section the rates of orbital evolution are conrputed for three planets. takine the fiux distribution or ΠΟ 209[58h and ILAT-P-2b from atmospheric models and for ID 1897323b from infrared observations.," In this section the rates of orbital evolution are computed for three planets, taking the flux distribution for HD 209458b and HAT-P-2b from atmospheric models and for HD 189733b from infrared observations." + Atinospheric+ iodels readily give the effective cluperature distribution T;g of the planet. cach patch of which radiates iuto a hemisphere.," Atmospheric models readily give the effective temperature distribution $T_{\rm{eff}}$ of the planet, each patch of which radiates into a hemisphere." + The thrust from a ch ix 6F=fol c., The thrust from a patch is ${\bf \delta F} = -f' \sigma T_{\rm{eff}}^4 {\bf \delta A} / c$ . + Tere 6A is the outward rormal vector to a patchoA of area dA. 6 is the Stefan- constant. aud the factor f’ accounts for he fact that the radiation is not all emitted normal to he surface: if it were. f/ would be 1.," Here ${\bf \delta A}$ is the outward normal vector to a patch of area $\delta A$, $\sigma $ is the Stefan-Boltzmann constant, and the factor $f'$ accounts for the fact that the radiation is not all emitted normal to the surface: if it were, $f'$ would be $1$." +" If treated as a rue blackbody. the radiation from a patch muaiformly ilunuinates the lemisphere. aud f/=3,"," If treated as a true blackbody, the radiation from a patch uniformly illuminates the hemisphere, and $f'=\onehalf$." + On the other iud. a limb-darkened atmosphere prefercutially cuits straight up rather than sidewavs. so. f/ is between i ancl L.," On the other hand, a limb-darkened atmosphere preferentially emits straight up rather than sideways, so $f'$ is between $\onehalf$ and $1$." + For iustauce. in à plane-parallel οταν atmosphere with intensity proportional to (p| 2/3). where µ is the cosine of the anele to the zenith (c.e.. 2)). this factor becomes fo=T.," For instance, in a plane-parallel gray atmosphere with intensity proportional to $(\mu+2/3)$ , where $\mu$ is the cosine of the angle to the zenith (e.g., \citealt{1978M}) ), this factor becomes $f'=4/7$." + Inteerating 6F over the surface of the planet eives the net force., Integrating ${\bf \delta F}$ over the surface of the planet gives the net force. + Finally. inteerating equations (1)) and (5)) over an orbit gives the secular effect of this force on the orbital elements.," Finally, integrating equations \ref{eq:dadt_md}) ) and \ref{eq:dedt}) ) over an orbit gives the secular effect of this force on the orbital elements." + For the planet IID209158b. ?.— carried out a detailed atmospheric circulation model aud ? ran radiative transfer calculations to predict phase curves.," For the planet HD209458b, \cite{2006CS} carried out a detailed atmospheric circulation model and \cite{2006F} ran radiative transfer calculations to predict phase curves." + For the equilibrimu-chemistiy case frou 7.. we compute the dipole componeut of the fux from the planet (based ou the lemispherc-averaged Tig as a fiction of orbital phase) to have an amplitude of L9«1075 ere sf. with ifs maxim shifted an augle 29° eastward.," For the equilibrium-chemistry case from \cite{2006F}, we compute the dipole component of the flux from the planet (based on the hemisphere-averaged $T_{\rm{eff}}$ as a function of orbital phase) to have an amplitude of $4.9 \times 10^{28}$ erg $^{-1}$, with its maximum shifted an angle $29^\circ$ eastward." + The timescale for somuinajor axis change is then 3.8«Lott ντ; for a fractional change of Aaja=0.3% over the age of the system. C, The timescale for semimajor axis change is then $3.8 \times 10^{11}$ yr; for a fractional change of $\Delta a / a = 0.3\%$ over the age of the system. ( +ALL age estimates are taken from ?..),All age estimates are taken from \citealt{2008TWH}. .) + The non-equilibrium chemistry case has a darker night side. so the dipole component is bigger. chhancing the Yarkovskv effect by about 20%," The non-equilibrium chemistry case has a darker night side, so the dipole component is bigger, enhancing the Yarkovsky effect by about $20\%$." + 2 have computed theoretical atinosphiere models for a ΡΟ of eccentric hot Jupiters., \cite{2007LL} have computed theoretical atmosphere models for a number of eccentric hot Jupiters. + Hore we analyze their computation for the eccentric planet ITAT-P-2h. choosing it because it(1) has a rather large eccentricity (0.5) which complements the analvsis dn and(2) is orbiting a bright star and will therefore receive detailed observatious.," Here we analyze their computation for the eccentric planet HAT-P-2b, choosing it because it(1) has a rather large eccentricity (0.5) which complements the analysis in \\ref{sec:calc} and(2) is orbiting a bright star and will therefore receive detailed observations." + See 7/— for the details of the model: we acditionallyassume is that there is no lamb darkening., See \cite{2007LL} for the details of the model; we additionallyassume is that there is no limb darkening. + The model was run for 19.5 orbits, The model was run for 19.5 orbits +subtracted.,subtracted. + Finally. the spectra were divided by the mean continu flux density to convert them iuto optical depth (since the line is optically thin. a simple division by the continu level suffices to convert from fiux density iuto optical depth).," Finally, the spectra were divided by the mean continuum flux density to convert them into optical depth (since the line is optically thin, a simple division by the continuum level suffices to convert from flux density into optical depth)." + CGAIRT observations of the2=0.2212 absorber were carried out on 19 October. 2000.," GMRT observations of the$z = 0.2212$ absorber were carried out on 19 October, 2000." + The standard 30station EX correlator. which gives a fixed ummber of 128 channels over a bandwidth which can be varie between 6I kIIz and 16 MIIz. was used as the backed.," The standard 30–station FX correlator, which gives a fixed number of 128 channels over a bandwidth which can be varied between 64 kHz and 16 MHz, was used as the backend." + A bandwidth of 1l MIIz was used for the observations. vieldiug a spectral resolutiou of ~2.0," A bandwidth of 1 MHz was used for the observations, yielding a spectral resolution of $\sim 2.0$." + L3CIIT ane οςσυ were used o calibrate the abxlute flux deusitv scale and svsteni xuidpass: uo phase calibrator was used since OI 363 is itself unresolved by t1ο QMBT., 3C147 and 3C286 were used to calibrate the absolute flux density scale and system bandpass; no phase calibrator was used since OI 363 is itself unresolved by the GMRT. + The toal on-source tinie was around four aid a half hours. with sixteen auteunuas oyeseut.," The total on-source time was around four and a half hours, with sixteen antennas present." + The data were analvsed in AIPS using standard srocedures., The data were analysed in AIPS using standard procedures. + The aalysis was fairly straightforward since here is ueelieible extended cussion in the OI 363 field., The analysis was fairly straightforward since there is negligible extended emission in the OI 363 field. + Coutimmin ciissiou was subtracted by fitting a iuear baseline to the U-V visibilities. using the AIPS ask UVLIN.," Continuum emission was subtracted by fitting a linear baseline to the U-V visibilities, using the AIPS task UVLIN." + The coutimmun subtracted data were then uapped nu all frequency channels aud. ai spectra extracted at the quasar location from the resulting threc-dimensional data cube., The continuum subtracted data were then mapped in all frequency channels and a spectrum extracted at the quasar location from the resulting three-dimensional data cube. + Spectra were also extracted at other locations in the cube to ensure that t10 data were not corrupted by interference., Spectra were also extracted at other locations in the cube to ensure that the data were not corrupted by interference. + A spectrum was also obtained by simply vector averaging the U-V visibilitics together. usiιο the ΑΠ task POSSAL," A spectrum was also obtained by simply vector averaging the U-V visibilities together, using the AIPS task POSSM." + Since the two methods are essentially equivalent for the preseut case of a point source at the telescope phase centre. we use the POSSAL specrun in the ciscussion below.," Since the two methods are essentially equivalent for the present case of a point source at the telescope phase centre, we use the POSSM spectrum in the discussion below." + The RAIS noise ou this spectrun ds ~2.5 mid. while the fix density of OI 363 was measured to be 2.25 Jy.," The RMS noise on this spectrum is $\sim 2.8$ mJy, while the flux density of OI 363 was measured to be 2.25 Jy." + Our experience with the GMBT indicates that the fiux density calibration is reliable to ~15. in this observing mode.," Our experience with the GMRT indicates that the flux density calibration is reliable to $\sim 15$, in this observing mode." + The spectra are shown in Fie., The spectra are shown in Fig. + 1 and Fig. 2. , \ref{fig:spectra} and Fig. \ref{fig:high}. . +Fig. ΤΑ) , Fig. \ref{fig:spectra}[ [ +presents the final Arecibo 3.125-MIIz (~0.395 resolution) spectrum in open squares.,A] presents the final Arecibo 3.125-MHz $\sim 0.395$ resolution) spectrum in open squares. + This has heen smoothed to a resolution of —2 NET compare it to the GAIRT spectrum. shown here as a thin solid Ime.," This has been smoothed to a resolution of $\sim 2$ to compare it to the GMRT spectrum, shown here as a thin solid line." + The agreement between the two is excellent., The agreement between the two is excellent. + The central narrow absorption compoucut can be seen to slightly asviunietric iu both spectra., The central narrow absorption component can be seen to be slightly asymmetric in both spectra. + Further. in additi to the deep narrow componcut. the 21-02 spectruii also has a shallow broad component. secu in both the Areci aud the CAIRT spectra.," Further, in addition to the deep narrow component, the 21-cm spectrum also has a shallow broad component, seen in both the Arecibo and the GMRT spectra." + This component cau be more clearly secu in the zooned-in plot of Fig. 1||, This component can be more clearly seen in the zoomed-in plot of Fig. \ref{fig:spectra}[ [ +D]. where the two spectra have again been superposed: again. the solid ine is the CAIRT spectiiun. while the open squares are he points ou the Arecibo 3.125-MITz spectimm.,"B], where the two spectra have again been superposed; again, the solid line is the GMRT spectrum, while the open squares are the points on the Arecibo 3.125-MHz spectrum." + Fig. ο , Fig. \ref{fig:spectra}[ [ +shows the Arecibo 3.125-MITIz spectrum at the origina resolution of ~0.1 (open squares): the thin solic ine is the multi-Caussian fit to the spectrmu which wil )o discussed. later.,C] shows the Arecibo 3.125-MHz spectrum at the original resolution of $\sim 0.4$ (open squares); the thin solid line is the multi-Gaussian fit to the spectrum which will be discussed later. + This spectra has au RAIS noise of ~0.00] iu optical depth per ~0.1 channel (1.0. ~2inJv)., This spectrum has an RMS noise of $\sim 0.001$ in optical depth per $\sim 0.4$ channel (i.e. $\sim 2$ mJy). + The deepest absorption occurs at a heliocentric requeucyv of 1163.075 MITz. ic. ata welioceutric redshift of :1.221250x0.000001.," The deepest absorption occurs at a heliocentric frequency of 1163.075 MHz, i.e. at a heliocentric redshift of $z = 0.221250 \pm 0.000001$." + For optically thin absorption by a homogenous III cloud. the 21-au optical depth 21 is related to the column density Nyy aud spin temperature T; by where fois the covering factor of the absorbing eas.," For optically thin absorption by a homogenous HI cloud, the 21-cm optical depth $\tau_{21}$ is related to the column density $\NHI$ and spin temperature $T_s$ by where $f$ is the covering factor of the absorbing gas." + OD 363 is a highly compact radio sotτος (Laneetal.2000.Stanehellinietal. 1997)) aud. as discussed in Cheugalur I&auekar (1999) and. Lane et al. (," OI 363 is a highly compact radio source \cite{lane00,stanghellini97}) ) and, as discussed in Chengalur Kanekar (1999) and Lane et al. (" +2000). the coveriug facor of the absorber is likely to be close to unity.,"2000), the covering factor of the absorber is likely to be close to unity." +" As lueitfioned m Sect. νι,"," As mentioned in Sect. \ref{intro}," + the earlier. lower seusitivitv CAIRT observations of Cheugalur Ranekar (1999) vielded a Inm1 spiu temperature 7;=112043:200 Is for the absorber.," the earlier, lower sensitivity GMRT observations of Chengalur Kanekar (1999) yielded a high spin temperature $T_s = 1120 \pm 200$ K for the absorber." + The deeper Arecibo observatious cau be used to obtain a niore precise estimate of the spin temperature of the svsteni., The deeper Arecibo observations can be used to obtain a more precise estimate of the spin temperature of the system. + The 3.125-MITIz Arecibo spectrui was used for this purpose as this. of all f16 Arecibo spectra. provides the optimum balance between resolution and sigual-to- ratio.," The 3.125-MHz Arecibo spectrum was used for this purpose as this, of all the Arecibo spectra, provides the optimum balance between resolution and signal-to-noise ratio." + The inteerated ootical depth of the spectiuu in Fig. 1|, The integrated optical depth of the spectrum in Fig. \ref{fig:spectra}[ [ +C] gives an IIT column denusitv of Nyy=O.ss9+0.007<1015Ty for the absorbing eas.,C] gives an HI column density of $\NHI = 0.889 \pm 0.007 \times 10^{18}~T_s$ for the absorbing gas. + The colui1 density obtained from the UST Liuau-o. profile is Nyy=7941.1107 (Bao&Turushek 1998))., The column density obtained from the HST $\alpha$ profile is $\NHI = 7.9 \pm 1.4 \times 10^{20}$ \cite{rao98}) ). + Comparing the two vields a spin temperature 7;=890c160 I. which agrees within the errors with the values obained carlicr x Lane et al. (," Comparing the two yields a spin temperature $T_s = 890 \pm 160$ K, which agrees within the errors with the values obtained earlier by Lane et al. (" +1998) aud Cheneahw Ίναiekar (1999).,1998) and Chengalur Kanekar (1999). + Iu the Calaxy. the neutral ISAL las no stable phase with a temperatire of ~1000 Ix. (although recent Arecibo observations indicate hat a significant fraction of neutral eas nav be in the τιstable phase with 500 Is Ty<5000. Ik: Teiles&Trelaud2000.Ποιος 2001)).," In the Galaxy, the neutral ISM has no stable phase with a temperature of $\sim 1000$ K (although recent Arecibo observations indicate that a significant fraction of neutral gas may be in the unstable phase with $500$ K $O3kpe(e/G00 kin/s)).," Although the use of $(\ref{eq:term})$ assumes a constant $\Gamma_0$ along a streamline, the distance travelled by the wind over the time-scale of change of $\Gamma_0$ $\sim 10 Myr$ ) is large $l \ge 6.3 \, {\rm kpc} (v/600 \, {\rm km/s})$ )." + Iu other words. the tiue-scale for the wind to reach a considerable wight above the disk is comparable to the time-scale of evolution iu Do. and therefore we can use our fornalisui ο estimate he terminal speed with the evolution in Dy.," In other words, the time-scale for the wind to reach a considerable height above the disk is comparable to the time-scale of evolution in $\Gamma_0$, and therefore we can use our formalism to estimate the terminal speed with the evolution in $\Gamma_0$." +" Figure 3 shows the evoution of οςfe, with time for >=0 (solic luec) and :=7 (dotted lino).", Figure 3 shows the evolution of $v_\infty /v_c$ with time for $z=0$ (solid line) and $z=7$ (dotted line). +" The curves show that tjo wind spec decreases rapidly after ~10 Abvr. aud fmat eyBe, its value beime smaller for compact galaxies."," The curves show that the wind speed decreases rapidly after $\sim 10$ Myr, and that $v_\infty \le 3 v_c$, its value being smaller for compact galaxies." + This result can be compared with the observed range of maxiunni wind speed., This result can be compared with the observed range of maximum wind speed. + ALartin (2005) fouud that he maxiuun speed of clouds embedded iu outflowiue eas ranges between 2 30.. aud Rupke (2005) found a rauge of [0.673|Co ," Martin (2005) found that the maximum speed of clouds embedded in outflowing gas ranges between $2\hbox{--}3 \, v_c$ , and Rupke (2005) found a range of $[0.67\hbox{--}3] \, v_c$." +We show this ranec with two dashed lines in Fie 3., We show this range with two dashed lines in Fig 3. + Figure 3 also shows that the wind speed is somewhat smaller at high redshift., Figure 3 also shows that the wind speed is somewhat smaller at high redshift. + The reason is that galactic mass for a given c. is παπα]: at hieh redshift. but cxAU2(1|i)1 and the mass effect outweighs the redsüft effect.," The reason is that galactic mass for a given $v_c$ is smaller at high redshift, but $c\propto M^{-0.2} \, (1+z)^{-1}$, and the mass effect outweighs the redshift effect." + However. the variation of οςes with redshift is expected to be very μα] iu this model. ich snaller than those caused by other parieters. such as the IME.," However, the variation of $v_\infty/v_c$ with redshift is expected to be very small in this model, much smaller than those caused by other parameters, such as the IMF." +" Figure 3 shows that dus-driven winds are likely to inve a tenniual speed ~23e,. for à combination of reasons that involve stellar plysics aud the relation )etwoeen disk :uid halo paraueters,"," Figure 3 shows that dust-driven winds are likely to have a terminal speed $\sim 2\hbox{--}3 v_c$, for a combination of reasons that involve stellar physics and the relation between disk and halo parameters." + It is interesting that his result coincides with observations. since there is ιο free parameter in our caletlation.," It is interesting that this result coincides with observations, since there is no free parameter in our calculation." + The strength of our approach les in the fact that the terminal speed culated usiic the Bernoulli function is iudepeudenut of he streamline used by the gas. streamlines do end to infinity. which is our sje asstuuption.," The strength of our approach lies in the fact that the terminal speed calculated using the Bernoulli function is independent of the streamline used by the gas, streamlines do extend to infinity, which is our basic assumption." +" Below discuss a ft""v implications.", Below we discuss a few implications. + Iu the sceneuio of energv driven winds. the ICAL is believed. to le| enriched. by winds from dyurf galaxies. since they were nuuerous in the early universe (Silk. Wrse ShickIs 1987: Nath Treutham 1997: Ferrara. Pettini Sheickiuov 2000: Con Bryan 2001: Madan. Ferrara. Rees 2001: Aguirre 2001).," In the scenario of energy driven winds, the IGM is believed to be enriched by winds from dwarf galaxies, since they were numerous in the early universe (Silk, Wyse Shields 1987; Nath Trentham 1997; Ferrara, Pettini Shchekinov 2000; Cen Bryan 2001; Madau, Ferrara, Rees 2001; Aguirre 2001)." + Towever. in the case of dustariven winds. the miportauce of low-nmiass ealaxies iun ICM curichiuent diiinishes because esX0s.," However, in the case of dust-driven winds, the importance of low-mass galaxies in IGM enrichment diminishes because $v_\infty \propto v_c$." + Qur calculatious here show that the wind speed depends strouely on the time elapsed. after a starburst. or more eoncrally on fje star formation history aud parameters.," Our calculations here show that the wind speed depends strongly on the time elapsed after a starburst, or more generally on the star formation history and parameters." + ]t is believedhat the IMP is weighted towards massive stars at high xcvdshift (c.g.. Schneider Oninkai 2010). iu which case the (vind speed likely increases with redshift.," It is believedthat the IMF is weighted towards massive stars at high redshift (e.g., Schneider Omukai 2010), in which case the wind speed likely increases with redshift." + Iu this case tie contribution of dwarfgalaxies at high redshift may still be important., In this case the contribution of dwarfgalaxies at high redshift may still be important. +the characteristics of sslars. which will be discussed briefly in this paper. caution us to keep that possibilitw in mind.,"the characteristics of stars, which will be discussed briefly in this paper, caution us to keep that possibility in mind." + In addition to their chemical signature. (his class of stars is also characterized by a limited range in spectral types (early F ancl A type stars) (Solanoetal.2001).," In addition to their chemical signature, this class of stars is also characterized by a limited range in spectral types (early F and A type stars) \citep{Solanoetal01}." +. For a long time they were also thought to be strictly limited to voung stars (ZAMS or but the evidence now points to ages ranging from the ZAMS to the TANS 2002).," For a long time they were also thought to be strictly limited to young stars (ZAMS or pre-main-sequence) but the evidence now points to ages ranging from the ZAMS to the TAMS \citep{IlievBarzova95,Ilievetal02}." +. A majority of sslars are pulsating stars in the general class of Seuli--(wpe variable stars 1999)., A majority of stars are pulsating stars in the general class of -type variable stars \citep{Bohlenderetal99}. +. Such pulsations indicate that the superficial helium. abundance is roughly solar. or higher.," Such pulsations indicate that the superficial helium abundance is roughly solar, or higher." + As a much larger fraction of sstous than other Α-ίνρο stars are vvariables (Paunzenοἱal.1998).. it suggests that the pphenomenon itself has an effect in their pulsational behavior.," As a much larger fraction of stars than other A-type stars are variables \citep{Paunzenetal98}, it suggests that the phenomenon itself has an effect in their pulsational behavior." + Finally. manv sstars show signs of circumstellar matter (Holwegeretal.1999).," Finally, many stars show signs of circumstellar matter \citep{Holwegeretal99}." +. several ideas have been put lorward (o account for the pphenomenon., Several ideas have been put forward to account for the phenomenon. + The very peculiar abundances of sslars cannot be reconciled with the diffusion models that have been so successful lor other chemically peculiar A stars such as the FinAm stars (Richeretal.2000)., The very peculiar abundances of stars cannot be reconciled with the diffusion models that have been so successful for other chemically peculiar A stars such as the FmAm stars \citep{RMT00}. +. Nevertheless. as dilfusion is an important process in A stars. il is an important feature of two leading models proposed for sslars. (he diffision/mass loss model (Michaud:audCharland1986:Charbonneau1993) and the acceretion/cdiffusion model (Vennand.Lambert1990;Charbonneau1991).," Nevertheless, as diffusion is an important process in A stars, it is an important feature of two leading models proposed for stars, the diffusion/mass loss model \citep{MichaudCharland86,C93} and the accretion/diffusion model \citep{VL90,C91}." +. A third hypothesis calls lor a binary merging to account for some Traction of sstars (Andrievsky1900), A third hypothesis calls for a binary merging to account for some fraction of stars \citep{Andrievsky97}. +", It also has been shown that a small number of sslars were in [act unidentified binaries lor which the combined spectra were mistakenly interpreted as metal-poor (Farageianaetal.2001).", It also has been shown that a small number of stars were in fact unidentified binaries for which the combined spectra were mistakenly interpreted as metal-poor \citep{Faraggianaetal01}. +. It is doubtful however that a significant fraction of sslars suller [rom this problem., It is doubtful however that a significant fraction of stars suffer from this problem. + The pro and cons of these hypothesis have been nicely summarized in Solanoetal.(2001)., The pro and cons of these hypothesis have been nicely summarized in \citet{Solanoetal01}. +. It is clear that at this moment no model can adequately reproduce all (he observed properties of sslars., It is clear that at this moment no model can adequately reproduce all the observed properties of stars. + The first challenge [acing the models is their inability to produce both: voung and old sslars., The first challenge facing the models is their inability to produce both young and old stars. + The cdiffusion/mass loss aud binary. merging models can only vield old sslars., The diffusion/mass loss and binary merging models can only yield old stars. + On the other hand the accretion/cdiffusion model as it stands now can only occur in, On the other hand the accretion/diffusion model as it stands now can only occur in +we see that (Vv)/p is unchanged. following the motion and C is conserved.,we see that $(\nabla \times {\bf v})/ \rho$ is unchanged following the motion and $\mathcal{C}$ is conserved. + The non-axisvmumetric exact solutions present a more demanding test. problem since the Toy Star changes shape throughout the oscillation., The non-axisymmetric exact solutions present a more demanding test problem since the Toy Star changes shape throughout the oscillation. + This changing shape must then be followed through the expansion and contraction. phases. maintaining the free boundary.," This changing shape must then be followed through the expansion and contraction phases, maintaining the free boundary." + The non-axisvmmetric Tov Star therefore presents an ideal problem for numerical codes designed to simulate astrophysical svstems with changing &cometries and [ree boundaries., The non-axisymmetric Toy Star therefore presents an ideal problem for numerical codes designed to simulate astrophysical systems with changing geometries and free boundaries. + The simulation is set up by again perturbing the equilibrium. solution. in this case with a velocity. of the form (28))-(29)) “Phe four parameters Vii.14s.τοι ancl Yos specifv the amplitude ancl geometry of the perturbation.," The simulation is set up by again perturbing the equilibrium solution, in this case with a velocity of the form \ref{eq:vxnonaxi}) \ref{eq:vynonaxi}) ) The four parameters $V_{11}, V_{12}, V_{21}$ and $V_{22}$ specify the amplitude and geometry of the perturbation." +the V-baud bolometric correetious (BC-) from several theoretical. etupirical. and. seui-eimpirical sources aud find that with a cousistent choice for the solar BC. all the sources agree to within 0.05 mag aloug a typical globular cluster isochrone.,"the $V$ -band bolometric corrections $(BC_V)$ from several theoretical, empirical, and semi-empirical sources and find that with a consistent choice for the solar $BC_V$, all the sources agree to within 0.05 mag along a typical globular cluster isochrone." + Our own analysis of different (BC)-) from different sources σοι this result., Our own analysis of different $(BC_V)$ from different sources confirm this result. + Therefore. for BC). we draw the uncertainty term [rom the uniform distribution +0.05 mae.," Therefore, for $BC_V$, we draw the uncertainty term from the uniform distribution $\pm$ 0.05 mag." + To explore uncertainties in παμεRy... we couducted a Monte Carlo simulation with 1120 incepenuclent sets of raudomly clioseu stellar evolution parameters. [rom each set generating Luminosity [unctious at several different metallicities. ages. aucl assumed initial mass functions.," To explore uncertainties in and, we conducted a Monte Carlo simulation with 1120 independent sets of randomly chosen stellar evolution parameters, from each set generating luminosity functions at several different metallicities, ages, and assumed initial mass functions." + Generating luminosity fuuctious requires a series of theoretical evolutionary tracks calculated for stars over a range of masses., Generating luminosity functions requires a series of theoretical evolutionary tracks calculated for stars over a range of masses. + In tliis work we use tlie nine masses Af = 0.55. 0.63. 0.70. 0.75. 0.80. 0.85. 0.90. 0.95. and 1.00 AZ..," In this work we use the nine masses $M$ = 0.55, 0.63, 0.70, 0.75, 0.80, 0.85, 0.90, 0.95, and 1.00 $M_\odot$." + Each evolutionary track is calculated [rom the zero-age tail sequence to the tip of the RGB., Each evolutionary track is calculated from the zero-age main sequence to the tip of the RGB. + The range of masses was chosen to allow reliable luminosity functions lor ages between 10 aud 20 Cigayears., The range of masses was chosen to allow reliable luminosity functions for ages between 10 and 20 Gigayears. + From a set of evolutionary tracks at appropriate masses oue cau Calculate theoretical isochrones ata variety of ages., From a set of evolutionary tracks at appropriate masses one can calculate theoretical isochrones at a variety of ages. + We use a moclilied version of the isochrone generating program used to coustruct the Revised Yale Isochroues (Green.Demarque&Wine1987)., We use a modified version of the isochrone generating program used to construct the Revised Yale Isochrones \citep{ryi}. +. The program uses the inethod of equivalent evolutionary points. locating in each evolutionary track a set of points delined in terms ¢ ‘the central belitum abuudance (ou the main sequence) or the helium core mass (on the RGB) and i-—terpolating among[n] them to generate the isoclroues.," The program uses the method of equivalent evolutionary points, locating in each evolutionary track a set of points defined in terms of the central helium abundance (on the main sequence) or the helium core mass (on the RGB) and interpolating among them to generate the isochrones." + From these isoclrones. V-baud Iuniuosity [uuctions are caleulated.," From these isochrones, $V$ -band luminosity functions are calculated." + Uncertainty in the bolometric correction is iucorporated at a later stage of the analysis., Uncertainty in the bolometric correction is incorporated at a later stage of the analysis. + Luminosity funetious are normalized to 1000 stars on tlie zero-age malin secqueuce aud use a bin size of 0.01 mag., Luminosity functions are normalized to 1000 stars on the zero-age main sequence and use a bin size of 0.04 mag. + Generating a lumiuosity function requires au initial mass fuuction (IME) describing the relative uumber[us of stars as a Duuction of mass created. during the formation of the cluster., Generating a luminosity function requires an initial mass function (IMF) describing the relative number of stars as a function of mass created during the formation of the cluster. +" The IMF is eenerally assumed to have the form of a simple declining power law. £A)xAL""."," The IMF is generally assumed to have the form of a simple declining power law, $\xi(M) +\propto M^{-\alpha}$." + In the classic studs. Salpeter(1955). found that observational data suggested anu initial mass [function with exponent a = 2.35: more recent studies. though. find evidence that the initial mass [function [or elobular clusters varies from cluster to cluster aud that metal-poor clusters in general have less steep IMFs than metal-rich clusters.," In the classic study, \citet{salpeter} found that observational data suggested an initial mass function with exponent $\alpha$ = 2.35; more recent studies, though, find evidence that the initial mass function for globular clusters varies from cluster to cluster and that metal-poor clusters in general have less steep IMFs than metal-rich clusters." + Because stars ou the RGB represent a narrow range of inasses. though. the luminosity function iu the red giaut region is less allected by the IME than are the main sequence aud subgiant brauch.," Because stars on the RGB represent a narrow range of masses, though, the luminosity function in the red giant region is less affected by the IMF than are the main sequence and subgiant branch." + Iu particular. we expect that aand wwill be unaffected by the INF.," In particular, we expect that and will be unaffected by the IMF." + This is tested by generating luminosity Dunctious using the, This is tested by generating luminosity functions using the +Combining the equations (AS--A9)) aud the photon momentum equations obtain the two motion constants / aud Q: where The above expression CAT1)) for the two motion coustanuts can be applied to calculations oft16 photon trajectory of auv poiut-like source with bulk motion iu νο metric.,Combining the equations \ref{eq:system6}- \ref{eq:system7}) ) and the photon momentum equations we obtain the two motion constants $l$ and $Q$: where The above expression \ref{eq:system9}) ) for the two motion constants can be applied to calculations of the photon trajectory of any point-like source with bulk motion in Kerr metric. +" If the source is static or has no bulk motion (io. J.j>0. 3,>0 aud jg» O0) this expression reduces to the simple oue given κ (1992)."," If the source is static or has no bulk motion (i.e. $\beta_{\varphi}\rightarrow 0$, $\beta_{r}\rightarrow 0$ and $\beta_{\theta}\rightarrow 0$ ), this expression reduces to the simple one given by \citet[]{kvp92}." +. For a flave above an accretion disk. where the flare bulk motion is observed in the locally uorotating faune to be predominantly in the z-directiou (say. upward or dowmward relative to the accretion ¢isk). them the z-direction velocity is proportional to the z-direction unit vector:," For a flare above an accretion disk, where the flare bulk motion is observed in the locally non-rotating frame to be predominantly in the $z$ -direction (say, upward or downward relative to the accretion disk), then the $z$ -direction velocity is proportional to the $z$ -direction unit vector:" +(b= 0.6).,$b = 0.6$ ). + Because ει>e2es the area with es/e1>e2/e1 is empty.," Because $e_1 > e_2 > e_3$ the area with $e_3/e_1 > +e_2/e_1$ is empty." + Note the different scales of the x and y-axis., Note the different scales of the $x$ and $y$ –axis. +" The vast majority of major branches has es/ei<0.1, which means that they are planar or close to planar."," The vast majority of major branches has $e_3/e_1 < 0.1$, which means that they are planar or close to planar." + Because of the nature of the branches some are even completely straight., Because of the nature of the branches some are even completely straight. +" In Figure 8,, we plot the fractional distributions of e2/e1 and e3/e1 for the three different halo groups b=0.60 (solid line), b=0.55 (dotted line), and b=0.50 (dashed line)."," In Figure \ref{shapes_histo}, , we plot the fractional distributions of $e_2/e_1$ and $e_3/e_1$ for the three different halo groups $b=0.60$ (solid line), $b=0.55$ (dotted line), and $b=0.50$ (dashed line)." +" As already seen in Figure 7,, just a few percent of branches have es/€1>0.1."," As already seen in Figure \ref{shapes_total}, just a few percent of branches have $e_3/e_1 > +0.1$." + The distribution of e2/ei is broader., The distribution of $e_2/e_1$ is broader. +" However, only a few percent of the branches have e2/ei>0.35."," However, only a few percent of the branches have $e_2/e_1 > 0.35$." + es/e170.1 and e2/e1<0.35 corresponds to straight or slightly curved filaments., $e_3/e_1 \approx 0.1$ and $e_2/e_1 \leq 0.35$ corresponds to straight or slightly curved filaments. +" In addition to comparing the fractional distributions of e2/e1 and es/ei, Figure 8 also compares halo groups for different linking lengths b."," In addition to comparing the fractional distributions of $e_2/e_1$ and $e_3/e_1$, Figure \ref{shapes_histo} also compares halo groups for different linking lengths $b$." + There are no discernable differences between the samples., There are no discernable differences between the samples. +" The mean values of these quantities for the different samples ε»/ει=0.1716, 0.1734, and 0.1736 and es/ei=0.0432, 0.0437, and 0.0428 for b—0.60, 0.55, and 0.50, respectively."," The mean values of these quantities for the different samples $\overline{e_2/e_1} = 0.1716$, 0.1734, and 0.1736 and $\overline{e_3/e_1} = +0.0432$, 0.0437, and 0.0428 for $b = 0.60$, 0.55, and 0.50, respectively." +" In a sense, the results obtained for the shapes of branches are not very surprising."," In a sense, the results obtained for the shapes of branches are not very surprising." + Haloes that comprise the largest object at b—0.50 are also contained in the largest object at larger b., Haloes that comprise the largest object at $b=0.50$ are also contained in the largest object at larger $b$. +" This statement is not trivial, though."," This statement is trivial, though." +" It is possible to have a distribution where the haloes in the largest object at b—0.50, say, are not contained in the one at b—0.55.9.."," It is possible to have a distribution where the haloes in the largest object at $b=0.50$, say, are not contained in the one at $b=0.55$." +" Given the results so far, LSS thus can be understood as being composed of a fixed set of building blocks, with different sizes and shapes, very much like a set of cosmic Lego bricks."," Given the results so far, LSS thus can be understood as being composed of a fixed set of building blocks, with different sizes and shapes, very much like a set of cosmic Lego bricks." +" Depending on how one chooses the linking length, the largest object is constructed by picking the same pieces, albeit in different numbers."," Depending on how one chooses the linking length, the largest object is constructed by picking the same pieces, albeit in different numbers." + Figure 9 shows the distributions of spatial extents of major branches., Figure \ref{histo_length} shows the distributions of spatial extents of major branches. +" By spatial extent we here mean for each major branch, where (min, Ymin, Zmin) and (Zmaz, Umaz; Zmaz) are the minimum and maximum coordinates of the major branch, respectively."," By spatial extent we here mean for each major branch, where $x_{min}$, $y_{min}$, $z_{min}$ ) and $x_{max}$, $y_{max}$, $z_{max}$ ) are the minimum and maximum coordinates of the major branch, respectively." +" For perfectly straight branches { is the actual length of the branch, for a branch that forms two sides of a rectangle, { is the length of the diagonal, etc."," For perfectly straight branches $l$ is the actual length of the branch, for a branch that forms two sides of a rectangle, $l$ is the length of the diagonal, etc." + We also divided the major branches into two categories., We also divided the major branches into two categories. + The first category (interior) contains those branches that connect to two (or more) major branches at each end., The first category (interior) contains those branches that connect to two (or more) major branches at each end. + The second category (fringe) encompasses those that connect to two (or more) major branches only at one end., The second category (fringe) encompasses those that connect to two (or more) major branches only at one end. +" In other words, the latter branches have a loose end."," In other words, the latter branches have a loose end." +" For example, in Figure 2,, branches A, B, D, F, and G belong to the second category, whereas only branches C and E are category one branches."," For example, in Figure \ref{major_branches}, branches A, B, D, F, and G belong to the second category, whereas only branches C and E are category one branches." +" For the Millennium Run haloes, the numbers of branches in the two categories are roughly equal."," For the Millennium Run haloes, the numbers of branches in the two categories are roughly equal." +" Figure 9 shows the distributions of | for all (solid line), interior(dotted line), and fringe (dashed line) branches (b— 0.6)."," Figure \ref{histo_length} shows the distributions of $l$ for all (solid line), interior(dotted line), and fringe (dashed line) branches $b = 0.6$ )." + Fringe branches appear to be shorter than interior ones., Fringe branches appear to be shorter than interior ones. + Abouttwo thirds of the major branches have extents of up to |=10-! MMpc.," Abouttwo thirds of the major branches have extents of up to $l = 10\,h^{-1}$ Mpc." + The other third extends to, The other third extends to +The crucial question is whether the published. quantitative AD models are testable by using azimutlally averaged resulis a common technique in astrophysics.,"The crucial question is whether the published, quantitative AD models are testable by using azimuthally averaged results – a common technique in astrophysics." +" We have argued above that our choice of telescope beams was made to encompassonly the radii over which the AD models sav the magnetic field dominates. not the very extended ISM where the magnetic fields passing through cores may indeed ""randomly wander."," We have argued above that our choice of telescope beams was made to encompass the radii over which the AD models say the magnetic field dominates, not the very extended ISM where the magnetic fields passing through cores may indeed “randomly” wander." + The issue is whether it is valid to test these models. or whether observations can only serve as “input” to AD models taken io be correct.," The issue is whether it is valid to test these models, or whether observations can only serve as “input” to AD models taken to be correct." + ALT want to take the value of 2). measured at anv of (he lour GBT positions as (he best measure of the magnetic field strength in each envelope. since field (wisting in the envelopes could put the Ποια in the plane of the skv ancl hence viekl ως=0.," MT want to take the value of $B_{los}$ measured at any of the four GBT positions as the best measure of the magnetic field strength in each envelope, since field twisting in the envelopes could put the field in the plane of the sky and hence yield $B_{los} = 0$." + The implication of this MT argument is that wilh severely (wisting field lines. perhaps one GBT beam would pick up the ~correct™ Dj; (o compare with the core value.," The implication of this MT argument is that with severely twisting field lines, perhaps one GBT beam would pick up the “correct” $B_{los}$ to compare with the core value." + But MT can accommodate whatever value we measure al any. envelope position. even if il is zero al all four positions.," But MT can accommodate whatever value we measure at any envelope position, even if it is zero at all four positions." + In their view. this would just mean Chat the field that is largely or partly along the line of sight at the core position has just twisted to be in the plane of the skv αἱ all envelope positions.," In their view, this would just mean that the field that is largely or partly along the line of sight at the core position has just twisted to be in the plane of the sky at all envelope positions." + Thev are saving that the AD theory is not testable by observations., They are saying that the AD theory is not testable by observations. + To do this would require measuring the mass and flux in a 3D Παν tube as it wanders around the Milkv. Way galaxy. which will never be possible.," To do this would require measuring the mass and flux in a 3D flux tube as it wanders around the Milky Way galaxy, which will never be possible." + The complaint that we cid not take this point into consideration is another wav of saving the AD theory is not testable., The complaint that we did not take this point into consideration is another way of saying the AD theory is not testable. + Moreover. the AIT cartoon of four fragments in a twisted fIux tube does not seem to be generally borne out by observations of dense cores.," Moreover, the MT cartoon of four fragments in a twisted flux tube does not seem to be generally borne out by observations of dense cores." + Dense cores such as the four we have studied are often part of a dense filament mapped in CO or other tracers., Dense cores such as the four we have studied are often part of a dense filament mapped in CO or other tracers. + If such filaments were flux tubes. the magnetic field would be along the filament.," If such filaments were flux tubes, the magnetic field would be along the filament." + This is sometimes seen. especially for lower density filaments.," This is sometimes seen, especially for lower density filaments." + But the dense filaanents in Taurus and Perseus (where our four cores reside) are perpendicular to Cae plane-ol-skv field mapped by linearly polarized backeround starlight polarization (Goodmanetal.1990:Goldsmith2008).. suggesting that the filaments were lormed by motions along magnetic fields rather than being flux tubes.," But the dense filaments in Taurus and Perseus (where our four cores reside) are perpendicular to the plane-of-sky field mapped by linearly polarized background starlight polarization \citep{Getal00, Getal08}, suggesting that the filaments were formed by motions along magnetic fields rather than being flux tubes." +limits for the spectral class of the exciting star where we did not detect any Continuum emission.,limits for the spectral class of the exciting star where we did not detect any continuum emission. + For NGC 02210. the spectral index we calculate for the central peak of the source between 6.7 and 8.5 (112 indicates that it is optically thick at these frequencies.," For NGC 6334F, the spectral index we calculate for the central peak of the source between 6.7 and 8.5 GHz indicates that it is optically thick at these frequencies." + The equations we have used to calculate the region parameters assume that they are optically thin., The equations we have used to calculate the region parameters assume that they are optically thin. + The cllect of the violation of this assumption is to reduce our estimate of the ionizing Ες of the exciting star. ellectively making it a lower limit.," The effect of the violation of this assumption is to reduce our estimate of the ionizing flux of the exciting star, effectively making it a lower limit." + So for the region NCGCGO334E. which has a central core that appears to be optically thick from our spectral index caleulations. the stellar types we derive are a lower limit.," So for the region NGC6334F, which has a central core that appears to be optically thick from our spectral index calculations, the stellar types we derive are a lower limit." + However. the elect is small since Caume Mutel (1987). derive similar parameters from observations at 15 Cillz where the optical depth clleet will be much smaller.," However, the effect is small since Gaume Mutel \shortcite{Ga1987} derive similar parameters from observations at 15 GHz where the optical depth effect will be much smaller." + Norris (1993) observed. 15 sites of 6.7-Gllz methanol maser emission anc found that a large fraction of their sample of 6.7- and 12.2-Gllz methanol maser sources have a simple curved or linear spatial distribution., Norris \shortcite{No1993} observed 15 sites of 6.7-GHz methanol maser emission and found that a large fraction of their sample of 6.7- and 12.2-GHz methanol maser sources have a simple curved or linear spatial distribution. + Based on their spatial distribution. we can divide all 6.7-Gllz methanol maser sources into one of two classes: those with a simple linear or curved morphology (e.g. G339.88-1.26). and those with a more complex morphology (c.g. NGC 6334E).," Based on their spatial distribution, we can divide all 6.7-GHz methanol maser sources into one of two classes: those with a simple linear or curved morphology (e.g. G339.88-1.26), and those with a more complex morphology (e.g. NGC 6334F)." + We have radio continuum observations associated with only two sites of 6.7-Cillz methanol maser emission. one from each class of spatial morphology.," We have radio continuum observations associated with only two sites of 6.7-GHz methanol maser emission, one from each class of spatial morphology." + No models have been suggested to explain the methanol masers with complex spatial distributions., No models have been suggested to explain the methanol masers with complex spatial distributions. + Thev may represent a different evolutionary phase of the star formation process. as suggested by Forster Caswell (1989) for complex OLL and maser distributions.," They may represent a different evolutionary phase of the star formation process, as suggested by Forster Caswell \shortcite{Fo1989} for complex OH and maser distributions." + “Phree possibilities have been suggested to explain the curvecl/linear spatial morphology: shocks fronts. collimatedjets. and circumstellar dises (Norrisetal.1993:Norris 1995).," Three possibilities have been suggested to explain the curved/linear spatial morphology: shocks fronts, collimated jets, and circumstellar discs \cite{No1993,No1995}." +. For each ofthese. Norris (1995) make a specific prediction about the relationship between the methanol masers and the region.," For each of these, Norris \shortcite{No1995} make a specific prediction about the relationship between the methanol masers and the region." + Ifthe methanol masers form in the shocks at the interface between the ionized ancl molecular gas. then we expect to see them near the edge of regions. às is often observed for OIL.," If the methanol masers form in the shocks at the interface between the ionized and molecular gas, then we expect to see them near the edge of regions, as is often observed for OH." + LW the methanol masers form in jets or collimated outllows. then we would expect to observe them distributed radiallv to the region.," If the methanol masers form in jets or collimated outflows, then we would expect to observe them distributed radially to the region." + However. if the methanol masers form. in. the circumstellar clises of voung stars. then we would expect to observe them approximately coincident with the continuum peak for the region.," However, if the methanol masers form in the circumstellar discs of young stars, then we would expect to observe them approximately coincident with the continuum peak for the region." + 2055.120 is one of the strongest sources of both 6.7- and 12.2-Gllz methanol maser emission., G339.88-1.26 is one of the strongest sources of both 6.7- and 12.2-GHz methanol maser emission. + The only published high-resolution. infrared. observations of the C339.88-1.26 region (Testietal.1994) show a strong peak. which appears to be slightly ollset [rom the position we measure for the region.," The only published high-resolution infrared observations of the G339.88-1.26 region \cite{Te1994} show a strong peak, which appears to be slightly offset from the position we measure for the region." + The methanol maser emission has an approximately linear spatial distribution at both. 6.7 and 12.2 Gllz (Norrisetal.1988:Norris1993).. ancl we measure the masers to be olfset. slightly to the south-west of the peak in the continuum emission. but not significantly so.," The methanol maser emission has an approximately linear spatial distribution at both 6.7 and 12.2 GHz \cite{No1988,No1993}, and we measure the masers to be offset slightly to the south-west of the peak in the continuum emission, but not significantly so." + This is the position predicted for the 6.7-Gllz methanol masers if they occur in. circumstellar disces., This is the position predicted for the 6.7-GHz methanol masers if they occur in circumstellar discs. + Unlike the eeneral case observed for ΟΙ masers. methanol masers do not lie toward the edge of the region (Gaume&Alutel1987).. which seems to rule out the hypothesis that the masers emanate from shocked gas.," Unlike the general case observed for OH masers, methanol masers do not lie toward the edge of the region \cite{Ga1987}, which seems to rule out the hypothesis that the masers emanate from shocked gas." + Conceivably these observations are compatible with the masers ling in a hiehly collimated outllow. as they are racial to the region.," Conceivably these observations are compatible with the masers lying in a highly collimated outflow, as they are radial to the region." + However. the distribution of the masers is highly linear and they do not have a wide velocity range as might be expected if they emanated from a high-velocity outflow.," However, the distribution of the masers is highly linear and they do not have a wide velocity range as might be expected if they emanated from a high-velocity outflow." + Further. we also require a double-sided jet. which is almost never seen in stars.," Further, we also require a double-sided jet, which is almost never seen in stars." + From our radio continuum images we estimate that a D0.5 star is required to produce the observed. region., From our radio continuum images we estimate that a B0.5 star is required to produce the observed region. + However. as the model we used does not take into account the absorption of UV. photons by dust. this is a lower limit on the spectral class of the exciting star.," However, as the model we used does not take into account the absorption of UV photons by dust, this is a lower limit on the spectral class of the exciting star." + Norris (1995) used simple modelling to show that the observed spatial ancl velocity distribution of the masers is consistent with Ixeplerian. motion. but were unable to derive a mass for the star.," Norris \shortcite{No1995} used simple modelling to show that the observed spatial and velocity distribution of the masers is consistent with Keplerian motion, but were unable to derive a mass for the star." + We find that the position of the masers with respect to the parent region also supports the hypothesis that the masers lie in circumstellar clises., We find that the position of the masers with respect to the parent region also supports the hypothesis that the masers lie in circumstellar discs. + The NGC 6334 region is the most active known region of OD star formation in the Galaxy (Llarvey&Gatley1983)., The NGC 6334 region is the most active known region of OB star formation in the Galaxy \cite{Ha1983}. +. The region has been the subject of several surveys at raclio ane infrared. wavelengths (AleBreenctal.1979:ItodrtguezGregor 1950)..," The region has been the subject of several surveys at radio and infrared wavelengths \cite{Mc1979,Ro1982,Lo1986,St1989a}." +. These observations show six main sites of racio emission and a similar number of clusters of infrared emission., These observations show six main sites of radio emission and a similar number of clusters of infrared emission. + Unfortunately. the nomenclature of the region is rather confusing.," Unfortunately, the nomenclature of the region is rather confusing." + The main site of 6.7-Gllz methanol maser emission. called 351.42|0.64 in the methanol maser literature. is associatoc Lawith the racio source NGC 02210. which is designated NCC 6334-1 at infrared waveleneths.," The main site of 6.7-GHz methanol maser emission, called G351.42+0.64 in the methanol maser literature, is associated with the radio source NGC 6334F, which is designated NGC 6334-I at infrared wavelengths." +incorporated into dust in the inner envelope or they have been released [rom grains by erosion processes in the outer envelope.,incorporated into dust in the inner envelope or they have been released from grains by erosion processes in the outer envelope. +" In order to test the livpotliesis that metal isocyanikdles are formed by gas-phase chemistry we lave investigated the syuthesis of MgNC vla radiative association reactious betwee ancl the cyanopolyvues (Duubar&Petrie 2002):: +HC,N + (with i= 3.5.7.9). followed by clissociative recombination reactions of the form: +e MeNCΗ."," In order to test the hypothesis that metal isocyanides are formed by gas-phase chemistry we have investigated the synthesis of MgNC via radiative association reactions between $^+$ and the cyanopolyynes \citep{dun02}: + + (with $n = 3,5,7,9$ ), followed by dissociative recombination reactions of the form: + +." + shows tlie radial abundance (cna?) ol MgNC in both the presence and. absence of shells., shows the radial abundance $^{-3}$ ) of MgNC in both the presence and absence of shells. + Iu. these caleulations we adopt an initial Mg/H» abundauce ratio of the derived abundances aud column densities are directly. proportional to this atio., In these calculations we adopt an initial $_2$ abundance ratio of $^{-5}$ — the derived abundances and column densities are directly proportional to this ratio. +" The effect of te deusity-euhanced shells is. ouce again. to move the peak of the cistributiou outward and to concentrate the abundance to the peaks of the gas deusity. iu this case with almost equal peaks in the 15"" and 29"" shells."," The effect of the density-enhanced shells is, once again, to move the peak of the distribution outward and to concentrate the abundance to the peaks of the gas density, in this case with almost equal peaks in the $15''$ and $29''$ shells." +" The total MeC coliuuun density is calculated to be 5.7*10""oE cm7. compared to observed values of (0.93-5)xLOM em7 (Cuélinetal.1995:Highberger&Ziurys 2000)."," The total MgNC column density is calculated to be $\times10^{13}$ $^{-2}$, compared to observed values of $\times10^{13}$ $^{-2}$ \citep{gue95, hig03}." +.. The domuüuant formation reaction is with HCZN which. despite its smaller abundance has the largest radiative association rate coellicient (5))). 6.5010.ρου).P can? * (Dunbar&Petrie2002).," The dominant formation reaction is with $_7$ N which, despite its smaller abundance has the largest radiative association rate coefficient ), $\times10^{-9}(T/300)^{-0.47}$ $^3$ $^{-1}$ \citep{dun02}." +. [It was suggested by Dunbar&Petrie(2002) that the inclusion of into chemical models for IRC+10216 might significantly reduce the abuudauces of the species alle help to reconcile the discrepancy between modelled: aux observed columu densities., It was suggested by \citet{dun02} that the inclusion of into chemical models for IRC+10216 might significantly reduce the abundances of the species and help to reconcile the discrepancy between modelled and observed column densities. + lu the preseut model. we find that the reaction of Mg with HC3N. HCN aud oN results in only ~1% reduction in their caleulated colui cleusities. which does not siguificantlv improve the match with observatious.," In the present model, we find that the reaction of $^+$ with $_5$ N, $_7$ N and $_9$ N results in only $\sim1$ reduction in their calculated column densities, which does not significantly improve the match with observations." + However. the depletion ol these species is dependent (in a roughly linear fashion). ou the initial Me abundauce eimmplovec.," However, the depletion of these species is dependent (in a roughly linear fashion), on the initial Mg abundance employed." + Compared to previous chemical iuodels of IRC+10216 (e.g.Nejad&MillarLOST:Brown&Millar2003:Aeiincdezetal. 2008).. the model. preseuted. here is unique in the inclusion of density-enhanced shells of gas aud dust.," Compared to previous chemical models of IRC+10216 \citep[\eg][]{nej87,che93,mil00,bro03,agu08}, the model presented here is unique in the inclusion of density-enhanced shells of gas and dust." + The shells included. in our model have pliysical parameters similar to the cust shells observed by Mauron&Hugeius(2000)., The shells included in our model have physical parameters similar to the dust shells observed by \citet{mau00}. +. As expected. the modeled. ¢dummn deusities differ from those calculated in the moclel by lillaretal.(2000) (referred to hereafter as JHB).," As expected, the modeled column densities differ from those calculated in the model by \citet{mil00} (referred to hereafter as MHB)." + gives the column densities Or tese two models for 33 species for which observational column «densities have been »ublished., gives the column densities for these two models for 33 species for which observational column densities have been published. + To highlight the elfeet of the censiv-enliauced shes on the chemistry ancl o permit Comparison with models without dleusity-euhanuced shells. the column densities rom the present inodel with no ceusity-enhanced shes (referος {ο as NS). are also given.," To highlight the effect of the density-enhanced shells on the chemistry and to permit comparison with models without density-enhanced shells, the column densities from the present model with no density-enhanced shells (referred to as NS), are also given." + Te couun cenusities caleulated by the hree dierent models are generally iu goo agreement with observatious. especially giver he complex morphology of the source.," The column densities calculated by the three different models are generally in good agreement with observations, especially given the complex morphology of the source." + Notable differences between the JHB anc 6S models. include a substantia recluction in the HC4N/HCSN column cdeusity ratio in the new model due to the enhancer j»hotodissociatiou rate used lor HC34N (Crom he RATEO6 database)., Notable differences between the MHB and NS models include a substantial reduction in the $_3$ $_5$ N column density ratio in the new model due to the enhanced photodissociation rate used for $_3$ N (from the RATE06 database). + The C4 : abtuiclance, The $_3$ N abundance +at all times.,at all times. +" Thus, it is the dependence of z;,—, on T that determines how the mass accretion rate through the tail and at R;, changes with time."," Thus, it is the dependence of $x_{\tau=1}$ on $T$ that determines how the mass accretion rate through the tail and at $R_{in}$ changes with time." +" To find z,-1(T) we need to set the left hand side of equation to unity and solve the resultant equation for x as a (38))function of T'.", To find $x_{\tau=1}(T)$ we need to set the left hand side of equation \ref{eq:thick_sol}) ) to unity and solve the resultant equation for $x$ as a function of $T$. +" Since here weconsider the case of rj(x,T=0)>1 the relation between z,—; and T is provided with sufficient accuracy by solving the equation for a given initial density distribution in the disk."," Since here weconsider the case of $\tpar(x,T=0)\gg 1$ the relation between $x_{\tau=1}$ and $T$ is provided with sufficient accuracy by solving the equation for a given initial density distribution in the disk." +" For the Gaussian initial distribution in the form (33)) one finds that Whenever zo>>1 and e oobservations have ὃνrevealed that roughly LO% of the neutral gas is outside the thin disc. at distances of a few kiloparsces (?7.andreferencestherein)...," In star-forming disc galaxies, high-sensitivity observations have revealed that roughly $10\%$ of the neutral gas is outside the thin disc, at distances of a few kiloparsecs \citep[and references + therein]{Oosterloo07, Fraternali09}." + Phe origin of this extraplanar gas is not completely understood. but. it is now widely accepted that a laree fraction of is produced by supernova-»owered: bubbles which drive itneutral and: partly. ionised eas out of the galactic disc., The origin of this extra–planar gas is not completely understood but it is now widely accepted that a large fraction of it is produced by supernova-powered bubbles which drive neutral and partly ionised gas out of the galactic disc. + This gas is expected to travel hrough the halo and eventually to fall back to the disc in a ime-scale of ~SO00Myr. a mechanism which is known as galactic fountain (222)..," This gas is expected to travel through the halo and eventually to fall back to the disc in a time-scale of $\sim 80-100\Myr$, a mechanism which is known as galactic fountain \citep{Shapiro76,Bregman80,HouckB90}." +" Thus. in the regions within a ew kiloparsecs from the galactic disc a large number of ealactic-Lountain ""cold"" clouds (ives in the Milkv Way) should be travelling through the ubiquitous coronal gas ancl an interaction between the two phases is inevitable. with important consequences for the kinematics of both."," Thus, in the regions within a few kiloparsecs from the galactic disc a large number of galactic-fountain “cold” clouds s in the Milky Way) should be travelling through the ubiquitous coronal gas and an interaction between the two phases is inevitable, with important consequences for the kinematics of both." + Unfortunately. the kinematics of the coronal gas is virtually unconstrained: as mentioned. most of the evidence for this gas is indirect. and from the available data it," Unfortunately, the kinematics of the coronal gas is virtually unconstrained: as mentioned, most of the evidence for this gas is indirect, and from the available data it" +accretion column but also to adequately take into account the soft X-ray emission from the accretion area on the white dwarf.,accretion column but also to adequately take into account the soft X-ray emission from the accretion area on the white dwarf. + Depending on the field topology. the magnetic flux density. and the ram pressure of the free-falling accretion stream. coupling of the accretion plasma onto the field line trajectories can occur at different locations within the binary reference frame. and the infalling material may be channeled into one or more aceretion regions on the white dwarf surface.," Depending on the field topology, the magnetic flux density, and the ram pressure of the free-falling accretion stream, coupling of the accretion plasma onto the field line trajectories can occur at different locations within the binary reference frame, and the infalling material may be channeled into one or more accretion regions on the white dwarf surface." + Constraining the geometry and location of the active X-ray emitting accretion region(s) is important to understanding the origin of X-ray soft and hard components and their flux balance., Constraining the geometry and location of the active X-ray emitting accretion region(s) is important to understanding the origin of X-ray soft and hard components and their flux balance. + Prominent light curve features and phase-resolved spectral modeling may provide insight into the system geometry., Prominent light curve features and phase-resolved spectral modeling may provide insight into the system geometry. + Since it is crucial to know their orbital phasing. we convert the photometric into orbital phases as deseribed in Sect.," Since it is crucial to know their orbital phasing, we convert the photometric into orbital phases as described in Sect." + 3.1. and refer to these values in the following discussion. which is valid based on the assumption of synchronous rotation.," \ref{sec:optphot} and refer to these values in the following discussion, which is valid based on the assumption of synchronous rotation." + These features are for instance the two bright phases in our X-ray light curves and the minimum in-between., These features are for instance the two bright phases in our X-ray light curves and the minimum in-between. + The almost identical fits to the bright-phase spectra strongly indicate that the complete soft X-ray emission originates in one and the same aceretion region., The almost identical fits to the bright-phase spectra strongly indicate that the complete soft X-ray emission originates in one and the same accretion region. + The minima around gop»=0.26 and 0.81 (eu0.45 and 0.0) are pronounced mainly in the soft X-ray flux and may have their origin m a variety of mechanisms., The minima around $\varphi_\mathrm{orb} = 0.26$ and $0.81$ $\varphi_\mathrm{phot} = 0.45$ and $0.0$ ) are pronounced mainly in the soft X-ray flux and may have their origin in a variety of mechanisms. + These include a total or partial self-eclipse of the accretion region by the white dwarf. absorption in the accretion stream. or random mass-transfer variations on timescales of several hours.," These include a total or partial self-eclipse of the accretion region by the white dwarf, absorption in the accretion stream, or random mass-transfer variations on timescales of several hours." + An occultation of the white dwarf by the secondary star can be excluded. since it would be expected to occur at inferior conjunction.," An occultation of the white dwarf by the secondary star can be excluded, since it would be expected to occur at inferior conjunction." + A self-eclipse occurs when the accretion region passes behind the limb of the white dwarf because of the stellar rotation., A self-eclipse occurs when the accretion region passes behind the limb of the white dwarf because of the stellar rotation. + Such à feature was found in the X-ray / EUV light curves of several polars such as (?).. MT Dra(?).. VV Pup(?).," Such a feature was found in the X-ray / EUV light curves of several polars such as , MT Dra, VV Pup." +. The X-ray spectra during phases of self-eclipse are characterized by à vanishing emitting area. at least of the soft X-ray component. while the intrinsic absorption remains nearly constant.," The X-ray spectra during phases of self-eclipse are characterized by a vanishing emitting area, at least of the soft X-ray component, while the intrinsic absorption remains nearly constant." + At à system inclination of -70°(?).. a partial or total self-eclipse 1s plausible for AI Tri. if the colatitude of the accretion region is higher than 20°.," At a system inclination of $i\sim 70\degr$, a partial or total self-eclipse is plausible for AI Tri, if the colatitude of the accretion region is higher than $20\degr$." + Of the two phase intervals with reduced X-ray emission detected for AI Tri. the major faint phase exhibits the characteristic signatures of a self-eclipse. indicating that larger parts of the accretion region might disappear behind the horizon of the white dwarf between You=0.61 and φομ=0.86.," Of the two phase intervals with reduced X-ray emission detected for AI Tri, the major faint phase exhibits the characteristic signatures of a self-eclipse, indicating that larger parts of the accretion region might disappear behind the horizon of the white dwarf between $\varphi_\mathrm{orb} = 0.61$ and $\varphi_\mathrm{orb} + = 0.86$." + During the Yow=0.21—0.31 interval. the hardness ratio ts not higher than zero. and soft and hard X-radiation are emitted at a similar level.," During the $\varphi_\mathrm{orb} = 0.21 - 0.31$ interval, the hardness ratio is not higher than zero, and soft and hard X-radiation are emitted at a similar level." + An eclipse of the accretion region. hence. is less likely to explain this phase.," An eclipse of the accretion region, hence, is less likely to explain this phase." + When the accretion column crosses the line of sight. it can obscure parts of the accretion region.," When the accretion column crosses the line of sight, it can obscure parts of the accretion region." + This manifests itself as a sharp dip in the soft X-ray light curves with a typical duration of Ay~0.1 and as significantly enhanced absorption in the X-ray spectrum., This manifests itself as a sharp dip in the soft X-ray light curves with a typical duration of $\Delta\varphi \sim 0.1$ and as significantly enhanced absorption in the X-ray spectrum. + Light curves and spectra of ΑΙ Tri display this behavior during the soft X-ray minimum around Yew»=0.26., Light curves and spectra of AI Tri display this behavior during the soft X-ray minimum around $\varphi_\mathrm{orb} = 0.26$. + The partial-covering absorption term im the spectral models rises markedly: the hardness ratio of the light curves increases from HR~-0.8 to values around zero (Fig. 2))., The partial-covering absorption term in the spectral models rises markedly; the hardness ratio of the light curves increases from $\mathrm{HR} \sim -0.8$ to values around zero (Fig. \ref{fig:lc1044}) ). + The short ingress and egress times and the duration Ag=0.1 of the light curve dip are also consistent with an origin in stream absorption., The short ingress and egress times and the duration $\Delta\varphi = 0.1$ of the light curve dip are also consistent with an origin in stream absorption. + The energy-dependent dip due to stream absorption and the broad faint phase in the X-ray light curves of ΑΙ Tri resemble the properties of the high-field system AR UMa and the two-pole aceretor MT Dra(?)., The energy-dependent dip due to stream absorption and the broad faint phase in the X-ray light curves of AI Tri resemble the properties of the high-field system AR UMa and the two-pole accretor MT Dra. +. Their X-ray emission can be explained in the same way., Their X-ray emission can be explained in the same way. + The faint phase might alternatively be interpreted as stream absorption and the X-ray light curve dip at gon=0.26 as a self-eclipse of the accretion region., The faint phase might alternatively be interpreted as stream absorption and the X-ray light curve dip at $\varphi_\mathrm{orb} = 0.26$ as a self-eclipse of the accretion region. + In consequence. the optical and the UV light curves would then be anti-correlated with the X-ray data. which would disagree with the standard picture of aceretion in polars in which the maximum flux in the UV range is seen when the aceretion region passes the line of sight — just at the phase of the stream eclipse.," In consequence, the optical and the UV light curves would then be anti-correlated with the X-ray data, which would disagree with the standard picture of accretion in polars in which the maximum flux in the UV range is seen when the accretion region passes the line of sight – just at the phase of the stream eclipse." + If the ultraviolet light curve. however. were dominated by the revolution of an extended aceretion column. we may see its irradiated. surface in anti-phase with the soft X-ray emission.," If the ultraviolet light curve, however, were dominated by the revolution of an extended accretion column, we may see its irradiated surface in anti-phase with the soft X-ray emission." + describe an anti-correlation of optical and soft X-ray flux for V," describe an anti-correlation of optical and soft X-ray flux for ," + describe an anti-correlation of optical and soft X-ray flux for Vu," describe an anti-correlation of optical and soft X-ray flux for ," + describe an anti-correlation of optical and soft X-ray flux for Vul," describe an anti-correlation of optical and soft X-ray flux for ," + describe an anti-correlation of optical and soft X-ray flux for Vul.," describe an anti-correlation of optical and soft X-ray flux for ," + describe an anti-correlation of optical and soft X-ray flux for Vul..," describe an anti-correlation of optical and soft X-ray flux for ," +415°. it would be easier to detect anisotropic ealaxy-galaxy lensing.,"$\pm 15^\circ$, it would be easier to detect anisotropic galaxy-galaxy lensing." + Unfortunately. the situation is not that simple in either of these cases.," Unfortunately, the situation is not that simple in either of these cases." + Weak lensing of the bright centres may. make their resulting images either rounder or more elliptical than their intrinsic image shape (i.c. bottom panel of Figure 12).," Weak lensing of the bright centres may make their resulting images either rounder or more elliptical than their intrinsic image shape (i.e., bottom panel of Figure 12)." + Suppose that one chooses a minimum image ellipticity for the bright centres. and that the computation of (6) and ~(8) is performed. using only those bright centres with observed ellipticity (ensi276a- ," Suppose that one chooses a minimum image ellipticity for the bright centres, and that the computation of $\gamma^+ (\theta)$ and $\gamma^- (\theta)$ is performed using only those bright centres with observed ellipticity $\epsilon_{\rm light} > \epsilon_{\rm cut}$." +Some fraction of the bright centres whose intrinsic ellipticity truly exceeds εως will. by weak lensing by foreground. galaxies. have their resulting images made rounder than their intrinsic ellipticitv.," Some fraction of the bright centres whose intrinsic ellipticity truly exceeds $\epsilon_{\rm cut}$ will, by weak lensing by foreground galaxies, have their resulting images made rounder than their intrinsic ellipticity." + As a result. some bright centres. with intrinsic ellipticities that ave larger than cou will. in fact. be rejected on because their observed (post-Iensing) images have ellipticities smaller than cou.," As a result, some bright centres with intrinsic ellipticities that are larger than $\epsilon_{\rm cut}$ will, in fact, be rejected on because their observed (post-lensing) images have ellipticities smaller than $\epsilon_{\rm cut}$." + Phe number of such bright centres that are allected by this will vary with the magnitude of the shear that they experience., The number of such bright centres that are affected by this will vary with the magnitude of the shear that they experience. + In addition to changing the ellipticitv. weak lensing of the bright centres may rotate the orientations of their symmetry axes (i0. top panel of Figure 12).," In addition to changing the ellipticity, weak lensing of the bright centres may rotate the orientations of their symmetry axes (i.e., top panel of Figure 12)." + Lone simply ries to narrow one's analysis region relative to the svnimctry axes of the bright centres. a problem: will occur if the wright centres have been weakly lensed.," If one simply tries to narrow one's analysis region relative to the symmetry axes of the bright centres, a problem will occur if the bright centres have been weakly lensed." + Any rotation of he symmetry axes of the bright centres causes the analysis region that one truly desires (i0. the region that brackets he directions of the major and minor axes of the projected 1alo mass) to be rotated with respect to the analvsis region hat one must actually use in practice (i.e. the region hat brackets the directions of the observed. major. and minor axes of the image of the bright centre).," Any rotation of the symmetry axes of the bright centres causes the analysis region that one truly desires (i.e., the region that brackets the directions of the major and minor axes of the projected halo mass) to be rotated with respect to the analysis region that one must actually use in practice (i.e., the region that brackets the directions of the observed major and minor axes of the image of the bright centre)." + Therefore. narrowing the analysis region may actually increase the discrepancy between the observed function. 5.(ο)(8). and the function that one would. measure if the intrinsic symmetry axes of the bright centres were known.," Therefore, narrowing the analysis region may actually increase the discrepancy between the observed function, $\gamma^+ (\theta) / \gamma^- (\theta)$, and the function that one would measure if the intrinsic symmetry axes of the bright centres were known." +" In this section we adopt a fiducial elliptical lens with velocity dispersion c,=150 km truncation radius wy=100b| kpe. and projected. halo ellipticity ο=0.3. and we construct Monte Carlo simulations that are identical to the Monte. Carlo simulations in Section 6 (Le. the simulations used to obtain the central panel of Figure 14)."," In this section we adopt a fiducial elliptical lens with velocity dispersion $\sigma_v = 150$ km $^{-1}$, truncation radius $x_t = 100~h^{-1}$ kpc, and projected halo ellipticity $\epsilon = 0.3$, and we construct Monte Carlo simulations that are identical to the Monte Carlo simulations in Section 6 (i.e., the simulations used to obtain the central panel of Figure 14)." + Shown in Figure 16 is the elfect of narrowing the analysis region when computing the ratio of the mean tangential shears., Shown in Figure 16 is the effect of narrowing the analysis region when computing the ratio of the mean tangential shears. +" The left panel of Figure 16 shows the observed and intrinsic functions. |6)/5.(8). when all sources within E45"" of the svmmetry axes of the central. elliptical lens are used for the calculations."," The left panel of Figure 16 shows the observed and intrinsic functions, $\gamma+^ (\theta) / \gamma^- (\theta)$, when all sources within $\pm 45^\circ$ of the symmetry axes of the central, elliptical lens are used for the calculations." + The right. panel of Figure 16 shows the same functions as the left. panel. but here only sources that are within 15 of the svmmoetrv axes of the central. elliptical lens are used. for the calculations.," The right panel of Figure 16 shows the same functions as the left panel, but here only sources that are within $\pm 15^\circ$ of the symmetry axes of the central, elliptical lens are used for the calculations." + From the right panel of Figure 16. it is clear that narrowing the analysis region (ic.. using only sources that are very close to the symmetry axes) increases the degree of anisotropy in the galaxv-galaxy lensine signal.," From the right panel of Figure 16, it is clear that narrowing the analysis region (i.e., using only sources that are very close to the symmetry axes) increases the degree of anisotropy in the galaxy-galaxy lensing signal." + However. narrowing the analvsis region. also increases the cisparity between the observed function. (0)(6). and the function that would be measured if the intrinsic svmimetry axes of the central. elliptical lenses were known.," However, narrowing the analysis region also increases the disparity between the observed function, $\gamma^+ (\theta) / +\gamma^- (\theta)$, and the function that would be measured if the intrinsic symmetry axes of the central, elliptical lenses were known." + Figure 17 shows the effect. of rejecting bright centres whose images (post-lensing) are very round., Figure 17 shows the effect of rejecting bright centres whose images (post-lensing) are very round. + ALL sources within +45° of the symmetry axes of the central. elliptical lens are used in the calculation.," All sources within $\pm 45^\circ$ of the symmetry axes of the central, elliptical lens are used in the calculation." + Here circles show the observed. function. «(8)/*5(06). and. crosses show the function that one would obtain if the intrinsic symmetry axes of the central elliptical lens were known.," Here circles show the observed function, $\gamma^+ (\theta) / +\gamma^- (\theta)$, and crosses show the function that one would obtain if the intrinsic symmetry axes of the central elliptical lens were known." + In the case of the circles ancl crosses. no constraint on the ellipticity of the lens image is imposed.," In the case of the circles and crosses, no constraint on the ellipticity of the lens image is imposed." + “Triangles in Figure 17 show the observed function. 5.(8)/5.(8). where the ratio has been computed: using the observed. symmetry axes of central. elliptical galaxies whose images (post-lensine) have ellipticities. isn0.3.," Triangles in Figure 17 show the observed function, $\gamma^+ (\theta) / \gamma^- (\theta)$, where the ratio has been computed using the observed symmetry axes of central, elliptical galaxies whose images (post-lensing) have ellipticities $\epsilon_{\rm light} > 0.3$." + From this figure. then. rejection of lenses with image ellipticities ei0.3. increases the observed function. +(0)(6). only slightly.," From this figure, then, rejection of lenses with image ellipticities $\epsilon_{\rm light} < 0.3$ increases the observed function, $\gamma^+ (\theta) / \gamma^- (\theta)$, only slightly." + In particular. rejection of the lenses with the rouncdest images does. not allow one to recover the function that one would measure it the intrinsic svnunetry axes of the central. elliptical lenses were known.," In particular, rejection of the lenses with the roundest images does not allow one to recover the function that one would measure if the intrinsic symmetry axes of the central, elliptical lenses were known." + We have investigated the theory of galaxy-ealaxy lensing by non-spherical dark matter haloes. which should give rise to an anisotropy in the tangential shear experienced by distant source galaxies.," We have investigated the theory of galaxy-galaxy lensing by non-spherical dark matter haloes, which should give rise to an anisotropy in the tangential shear experienced by distant source galaxies." + I£ each distant. source is lensed by only one [oreeround elliptical lens. and if the observed svnimetry axes of the elliptical lens correspond to the intrinsic symmetry axes of its projected dark matter halo. one would expect the signature of anisotropic galaxy-ealaxy lensing to manifest as (8)5 over a wide range of angular scales.," If each distant source is lensed by only one foreground elliptical lens, and if the observed symmetry axes of the elliptical lens correspond to the intrinsic symmetry axes of its projected dark matter halo, one would expect the signature of anisotropic galaxy-galaxy lensing to manifest as $\gamma^+ (\theta) > \gamma^- (\theta)$ over a wide range of angular scales." + Here ~(80) is the angular dependence of the mean tangential shear experienced. by sources whose azimuthal coordinates, Here $\gamma^+ (\theta)$ is the angular dependence of the mean tangential shear experienced by sources whose azimuthal coordinates +remnant (Stranieroetal1997:Dominguez1999).,"remnant \citep{scl97,dom99}." +. For this reason. we considered only stars. distributed according to the Salpeters AIF with masses in (he range 0.12xm1.2 AL..," For this reason, we considered only stars distributed according to the Salpeter's MF with masses in the range $0.12 \leq m \leq 1.2$ $_\odot$." + Moreover. we assumed that the contribution to the low mass population due to the mentioned remnants is practically negligible.," Moreover, we assumed that the contribution to the low mass population due to the mentioned remnants is practically negligible." +" Indeed. according to the Salpeters ME. (he ratio between the number of remnants whose progenitor had a mass around 7», and the number of stars with mass around n is ⊳↔⊲↕∏↽≻↕↽≻∪⋟∖⊽↕∐≸≟⊔⋯↥⋟∖⊽∏≺∢∐↕↽≻↕⋅∪↖≺↽↔↴≼↲∐∐∪↕∷∖⊽≀↧↴↕⋅≼↲⊔↥∪⋟∖⊽≼↲↖≺↽↔↴↕∖↽↕∐≸≟↕⋅↕⋟∖⊽≼↲↥∪↕⋅≼↲∐↓∐≀↕↴∐↥⋟∖⇁∖∖⊽↕⊔⊔⊔≀↧↪∖⋱∖⊽∣∣∣⋅⊔∐↲∐⋅ ∐⋅∪∐↓⊔∐↲≼↲⊳∖⊽∐∐↓≀↧↴∥↲⋝∖⊽↕∐⊳↔��∏⋅≀↧↴∐↕≼↲↕⋅∪≼↲↥≀↧↴↥⊔≤∍≤∍∏≀↧↴∐≼⇂∐∪∐∐∐≸≟⋯↲∠≼↲↥≀↧↴↥⊔≤∍≤∍≤∍↕⋟⋅⋅∣∣∣↙⇂≼⋝⇀∖↕⋅⇄⋝≃ 9.5(m— 0.45). with m.20.45 M. because stars with lower mass remnants cannot be evolved in a IIubble time."," Indeed, according to the Salpeter's MF, the ratio between the number of remnants whose progenitor had a mass around $m_p$ and the number of stars with mass around $m$ is Supposing that such progenitors are those giving rise to remnants with mass $m$ , then, from the estimates in \cite{scl97} and \cite{dom99}, $m_p ($ $_\odot) \simeq 9.5 (m - 0.45)$ , with $m>0.45$ $_\odot$ because stars with lower mass remnants cannot be evolved in a Hubble time." + Thus. substituting in Eq. (C1)).," Thus, substituting in Eq. \ref{nremn}) )," + where my is the lowest mass a remnant can have at (he assumed cluster age., where $m_l$ is the lowest mass a remnant can have at the assumed cluster age. + From fitting the abovecited estimates. this lower limit turns out to be 0.59 M... ," From fitting the above–cited estimates, this lower limit turns out to be $m_l\sim +(0.45\log t_{gc} - 1.2)^{-3.1}+0.45 \sim 0.59$ $_\odot$." +One can see that the ratio in Eq. (C2)), One can see that the ratio in Eq. \ref{ratio}) ) + is monotonically decreasing lor m> m. hence (he maximun takes place for the lowest mass class we used in the model. i.e. i=0.71 M... for which JNτι~0.05.," is monotonically decreasing for $m\geq m_l$ , hence the maximum takes place for the lowest mass class we used in the model, i.e. $m=0.71$ $_\odot$, for which $N_{\rm remn}/N_{0.71} \sim 0.05$." + Since in our numerical representation τινο0.01 CN is the total number of particles). then in this class there should have been about 5xLO£N remnants.," Since in our numerical representation $N_{0.71}/N \simeq 0.01$ $N$ is the total number of particles), then in this class there should have been about $5\times 10^{-4}N$ remnants." +" Thus. bearing in mind that the less populated mass class contains ~2x10.75V. we can reasonably affirm that the ME we assumed [for the initial conditions was not substantiallv allected. by stellar evolution neither at the initialtime /,,. nor later during the simulation (because it lastsmuch shorter than /,.)."," Thus, bearing in mind that the less populated mass class contains $\sim 2\times 10^{-3}N$, we can reasonably affirm that the MF we assumed for the initial conditions was not substantially affected by stellar evolution neither at the initialtime $t_{gc}$ nor later during the simulation (because it lastsmuch shorter than $t_{gc}$ )." +This work was supported by the Natural Sciences and Engineering Research Council of Canada.,This work was supported by the Natural Sciences and Engineering Research Council of Canada. + We would like to thank the stall of the JCMT for. their assistance with the SCUBA observations., We would like to thank the staff of the JCMT for their assistance with the SCUBA observations. + WC would like to thank Vicki Barnard for assistance determining pre-upgracde calibration FCFs and Bernd Welerling for assistance. in determining problematic Toso fits., KC would like to thank Vicki Barnard for assistance determining pre-upgrade calibration FCFs and Bernd Weferling for assistance in determining problematic $\tau_{\mathrm{CSO}}$ fits. + We also wish to thank an anonvmous referee for constructive comments., We also wish to thank an anonymous referee for constructive comments. + The James Clerk Maxwell Telescope is operated. on behalf of the Particle Physies and Astronomy Research Council of the United Ixingdom. the Netherlands Organisation for Nunite SENResearch.ι ancl the National Research Council of anges," The James Clerk Maxwell Telescope is operated on behalf of the Particle Physics and Astronomy Research Council of the United Kingdom, the Netherlands Organisation for Scientific Research, and the National Research Council of Canada." + 2=2. O5. 0.045 Rauchetal.(1998);Weinberg(1997):Burles&Twvtler.(1998):Nirkmanal.(2003):BennettetSpergel(2007):IXomatsu(2009).. Fukugitaetal.," $z=2$ $\Omega_B$ $ 0.045$ \cite{Rauch98, Weinberg97, BurTyt98, Kirkman03, +Bennett03, Spergel07, Komatsu09}. \cite{Fukugita98}." +"(1998).. Cen&Ostriker(19993):Croftetal.(2001):DavéYoshikawa(2003):Borganietal.(2004):Cen&Ostriker(2006) 2=0. 50% 10?-10* (0.1—1 Tz10"" T<10 IRaxanond&Sinith(1977).. Z.. "," \cite{CenOst99, Croft01, Dave01, Yoshikawa03, Borgani04, CenOst06} $z=0$ $50~\%$ $10^5$ $10^7$ $0.1-1$ $>10^6$ $<10^6$ \cite{Raymond77}, $_\odot$ \cite{Snowden97, Mush00, +McCammon02, Galeazzi07, Henley08} " +"data are taken from ?;; ? provided and data; HST/NICMOS and WFPC2 data, which we have carefully dereddened, are from ?themselves.;; SCUBA data at um are taken from ?;; and SCUBA um and um) data are from ?;; VLA data are from ?..","data are taken from \citet{1998A&A...330...97S}; ; \citet{2006ApJ...641..801R} provided and data; /NICMOS and WFPC2 data, which we have carefully dereddened, are from \citet{2000ApJ...528..276M}; SCUBA data at $\mu$ m are taken from \citet{1993MNRAS.260..844H}; and SCUBA $\mu$ m and $\mu$ m) data are from \citet{1999A&A...341..667M}; VLA data are from \citet{1983ApJ...273..128B}." +" ? report X-ray observations by on February 2 and 6, 2001 and by on May 9 and 21, 2001 with a photon index Il=2 for the parsec-scale jet component."," \citet{2004ApJ...612..786E} report X-ray observations by on February 2 and 6, 2001 and by on May 9 and 21, 2001 with a photon index $\Gamma=2$ for the parsec-scale jet component." + Data from the NED are shown as non-constraining points gray) for comparison., Data from the NED are shown as non-constraining points ) for comparison. + The H.E.S.S. upper limit based on observations in 2004 with a live time exposure of hh is reported in red (?).., The H.E.S.S. upper limit based on observations in 2004 with a live time exposure of h is reported in red \citep{2005A&A...441..465A}. +" We should also point out that as AA harbors a strongly absorbing dust lane, and because this source is extremely close and well resolved, the X-ray data should be then taken only as upper limits."," We should also point out that as A harbors a strongly absorbing dust lane, and because this source is extremely close and well resolved, the X-ray data should be then taken only as upper limits." + The contribution of the nuclear jet might be contaminated by the accretion disk of the AGN and by the binaries that are resolved in this object., The contribution of the nuclear jet might be contaminated by the accretion disk of the AGN and by the X-ray binaries that are resolved in this object. + In this case we would have only poor constraints on the emission process., In this case we would have only poor constraints on the emission process. + We assume here that all our selected high energy data come from SSC processes., We assume here that all our selected high energy data come from SSC processes. +" Figure [/] shows the SED of AA applying the multi-blob model in two cases, (i) assuming that the y-ray peak observed by is inverse Compton radiation lines) or (ii) assuming it to be synchrotron lines)."," Figure \ref{fig:CenA_SED} shows the SED of A applying the multi-blob model in two cases, ) assuming that the $\gamma$ -ray peak observed by is inverse Compton radiation ) or ) assuming it to be synchrotron )." +" We should also point out that the previous study by ? reports a variability in soft y-ray of about 10 days, implying rp<2x 10cm (see Eq. (2))),"," We should also point out that the previous study by \citet{1998A&A...330...97S} reports a variability in soft $\gamma$ -ray of about 10 days, implying $r_b < 2 \times 10^{17}$ cm (see Eq. )," + which is well satisfied by our parameters., which is well satisfied by our parameters. +" Given the results of our model in the first scenario with an IC bump in soft y-rays lines), the SSC emission of the central region would definitely not provide a flux sufficiently high to be detectable at VHE (see column 6 in Table at least for a SSC emission dominated by the nucleus."," Given the results of our model in the first scenario with an IC bump in soft $\gamma$ -rays ), the SSC emission of the central region would definitely not provide a flux sufficiently high to be detectable at VHE (see column 6 in Table \ref{tab:param}) ), at least for a SSC emission dominated by the nucleus." + This [I).holds even in the case of huge variations of the nuclear emission., This holds even in the case of huge variations of the nuclear emission. +" This conclusion concurs with ?,, who do not expect SSC emission by the nucleus or by the base of the jet of AA, but do expect VHE emission that could be detectable by current imaging atmospheric Ceerenkov telescopes facilities in the case of an external inverse Compton emission process on the host galaxy photon field."," This conclusion concurs with \citet{2003ApJ...597..186S}, who do not expect SSC emission by the nucleus or by the base of the jet of A, but do expect VHE emission that could be detectable by current imaging atmospheric Čeerenkov telescopes facilities in the case of an external inverse Compton emission process on the host galaxy photon field." +" For ?,, many Fanaroff-Riley type I II) radiogalaxies like AA could be TeV sources for which the weak nuclear y-ray emission would be absorbed and re-processed by inverse Compton on the starlight photon field, thus generating an isotropic y-ray halo."," For \citet{2006MNRAS.371.1705S}, many Fanaroff-Riley type I I) radiogalaxies like A could be TeV sources for which the weak nuclear $\gamma$ -ray emission would be absorbed and re-processed by inverse Compton on the starlight photon field, thus generating an isotropic $\gamma$ -ray halo." +" In our model, the lack of simultaneous data prevents us from further constraining the synchrotron bump, which has here a higher flux density than the selected data since we are considering a state of high y-ray activity."," In our model, the lack of simultaneous data prevents us from further constraining the synchrotron bump, which has here a higher flux density than the selected data since we are considering a state of high $\gamma$ -ray activity." +" In solid (“on-blob” case) and short dashed-long dashed (“inter-blob” case) green lines in Fig. [7],"," In solid (“on-blob” case) and short dashed–long dashed (“inter-blob” case) green lines in Fig. \ref{fig:CenA_SED}," + we present a SED of CenAA assuming now that the soft y-ray peak is of synchrotron origin., we present a SED of A assuming now that the soft $\gamma$ -ray peak is of synchrotron origin. +" In this case (see column 7 in Table [I)), we expect a detection of the core of AA at VHE by the H.E.S.S. telescope array within hh of observation."," In this case (see column 7 in Table \ref{tab:param}) ), we expect a detection of the core of A at VHE by the H.E.S.S. telescope array within h of observation." +" It should be noted that ? also predicted the synchrotron bump to be in the soft y-ray range and the inverse Compton bump to peak around TTeV in the context of SSC models, which comforts our latter model."," It should be noted that \citet{2001ApJ...549L.173B} also predicted the synchrotron bump to be in the soft $\gamma$ -ray range and the inverse Compton bump to peak around TeV in the context of SSC models, which comforts our latter model." + 00521-—36 is a FSRQ object with an optical jet located at z=0.055 (?).., $-$ 36 is a FSRQ object with an optical jet located at $z=0.055$ \citep[][]{1985AJ.....90.2207K}. +" The central SMBH has a mass of 3.3x105 Mc (?). ?,,"," The central SMBH has a mass of $\sim$ $ \times 10^8 M_{\sun}$ \citep[][]{2005ApJ...631..762W}. \citet{1996ApJ...459..169P}," +" and more recently ?,, mention the absence of superluminal motions in its jet, contrary to the case of 2273, implying that the beaming effect is much less important and thus strengthening their result on the viewing angle."," and more recently \citet{2002AJ....124..652T}, mention the absence of superluminal motions in its jet, contrary to the case of 273, implying that the beaming effect is much less important and thus strengthening their result on the viewing angle." + Indeed the only constraint on the jet orientation comes from ? who deduce 0=~30°+6? from SSC models., Indeed the only constraint on the jet orientation comes from \citet{1996ApJ...459..169P} who deduce $\theta \simeq 30\degr \pm 6\degr$ from SSC models. +" We should also note that 00521—36 seems to oscillate between a Seyfert-like and a LLac state (e.g.?),, making this source difficult to interpret within a pure non-thermal scenario, especially since we are confronted with non-simultaneous data."," We should also note that $-$ 36 seems to oscillate between a Seyfert-like and a Lac state \citep[e.g.][]{1981A&A...103L...1U}, making this source difficult to interpret within a pure non-thermal scenario, especially since we are confronted with non-simultaneous data." +" BeppoSAX observed 00521—36 on October 10, 1998 (?) in Fig. 8p)."," SAX observed $-$ 36 on October 10, 1998 \citep{2002babs.conf...63G} in Fig. \ref{fig:pks0521_SED}) )," +" and the Swift/XRT measurements points) taken on May 26, 2005 were obtained through the Online Analysis"," and the /XRT measurements ) taken on May 26, 2005 were obtained through the Online Analysis." + The data points in gray are from the NED., The data points in gray are from the NED. + We report Τοο[!.in red the upper limit at 2σ obtained in hh from observations by CANGAROO between 1993 and 1996 (?).., We report in red the upper limit at $2\sigma$ obtained in h from observations by CANGAROO between 1993 and 1996 \citep{1998A&A...337...25R}. +" We further used the EGRET data between MMeV and MMeV from ? and taken between July 12, 1994 and August 01, 1994."," We further used the EGRET data between MeV and MeV from \citet{1999ApJS..123...79H} and taken between July 12, 1994 and August 01, 1994." +" The blazar 00521—36 is associated with the source JJ0524—3630 in the Second EGRET Catalog (?),, but during cycle 4 this source was found to lie outside the confidence contour of EGRET."," The blazar $-36$ is associated with the source $-$ 3630 in the Second EGRET Catalog \citep{1995ApJS..101..259T}, but during cycle 4 this source was found to lie outside the confidence contour of EGRET." +" However, like ?,, we assume in this work the identification with 00521—36 to be valid, which is also pointed out by ?.."," However, like \citet{2006A&A...453..829F}, we assume in this work the identification with $-$ 36 to be valid, which is also pointed out by \citet{2002ApJ...579..136T}." + Figure [B] presents the SED of PKS00521—36 with the anticipated VHE emission (see column 8 in Table [I] for the corresponding parameters)., Figure \ref{fig:pks0521_SED} presents the SED of $-$ 36 with the anticipated VHE emission (see column 8 in Table \ref{tab:param} for the corresponding parameters). +" It seems unlikely that the X-rays are due to inverse Compton radiation since the inverse Compton scattering would be with photons from the radio/optical contribution, which is not very variable, coming from an extended part of the jet."," It seems unlikely that the X-rays are due to inverse Compton radiation since the inverse Compton scattering would be with photons from the radio/optical contribution, which is not very variable, coming from an extended part of the jet." +" Since the X-rays show a high degree of variability, they most certainly come from a compact region and are of synchrotron origin."," Since the X-rays show a high degree of variability, they most certainly come from a compact region and are of synchrotron origin." +" If the X-ray emission is truly due to the synchrotron process, we predictthat this LLac object should be marginally detectable by the present H.E.S.S. telescope array, and easily detectable by III and by the next generation of Ceerenkov arrays, such as the CTA project,which will detect sources down to ~0.1% of the Crab flux."," If the X-ray emission is truly due to the synchrotron process, we predictthat this Lac object should be marginally detectable by the present H.E.S.S. telescope array, and easily detectable by II and by the next generation of Čeerenkov arrays, such as the CTA project,which will detect sources down to $\sim$ of the Crab flux." +" Furthermore, if 00521—36 remains undetected at VHE, a misidentification with the EGRET source should be considered."," Furthermore, if $-$ 36 remains undetected at VHE, a misidentification with the EGRET source should be considered." +Post-Asymptotie Giant Branch stars there. PAGB stars} are low mass stars evolving between the asymptotic giant branch CAGB) and the white dwarf cooling track.,"Post-Asymptotic Giant Branch stars (here, PAGB stars) are low mass stars evolving between the asymptotic giant branch (AGB) and the white dwarf cooling track." + PAGB stars evolve from stars with initial masses in the range 0.8. to SAZ.., PAGB stars evolve from stars with initial masses in the range $M_\odot$ to $M_\odot$. + Thanks to mass loss on the red giant branch and principally on the AGB. the PAGB stars are widely considered to have masses of about O.SA/~ or less.," Thanks to mass loss on the red giant branch and principally on the AGB, the PAGB stars are widely considered to have masses of about $0.8M_\odot$ or less." + Evolution from the cool AGB star to a hot star at the beginning of the white dwarf cooling track is rapid with times of 10.000 years thought to be representative.," Evolution from the cool AGB star to a hot star at the beginning of the white dwarf cooling track is rapid with times of 10,000 years thought to be representative." + Gas lost previously by the star is ionized by the hot PAGB central star to form a planetary nebula., Gas lost previously by the star is ionized by the hot PAGB central star to form a planetary nebula. + The dust component of the mass loss may be detected as an infrared excess., The dust component of the mass loss may be detected as an infrared excess. + PAGB stars have been reviewed by Kwok (1993) and Van Winckel (2003)., PAGB stars have been reviewed by Kwok (1993) and Van Winckel (2003). + 0.2 em One reason for great interest in PAGB stars is that they have the potential to provide observational constraints. particularly through studies of their chemical compositions. on the complex mix of evolutionary processes — nucleosynthesis. mixing and mass loss — occurring on the AGB.," 0.2 cm One reason for great interest in PAGB stars is that they have the potential to provide observational constraints, particularly through studies of their chemical compositions, on the complex mix of evolutionary processes – nucleosynthesis, mixing and mass loss – occurring on the AGB." + Interpretation of these constraints for field PAGB stars is compromised in part because the composition and mass of the main sequence. red giant. and AGB progenitor are not directly known.," Interpretation of these constraints for field PAGB stars is compromised in part because the composition and mass of the main sequence, red giant, and AGB progenitor are not directly known." + Such compromises are essentially eliminated by finding a PAGB star as a member of an open or globular cluster., Such compromises are essentially eliminated by finding a PAGB star as a member of an open or globular cluster. + 0.2 em In this paper. we report on an abundance analysis of the A-type PAGB star discovered by Siegel Bond (2009. in preparation) in the globular cluster M79.," 0.2 cm In this paper, we report on an abundance analysis of the A-type PAGB star discovered by Siegel Bond (2009, in preparation) in the globular cluster M79." + The location of the star in the colour-magnitude diagram is shown in Figure |.., The location of the star in the colour-magnitude diagram is shown in Figure \ref{f_m79_color}. + The initial mass of this star must have been slightly in excess of the mass of stars now at the main sequence turn-off. say. AJ2O.SAL..," The initial mass of this star must have been slightly in excess of the mass of stars now at the main sequence turn-off, say, $M \simeq 0.8M_\odot$." + The star's composition may be referenced to that of the cluster's red giants. stars for which abundance analyses have been reported.," The star's composition may be referenced to that of the cluster's red giants, stars for which abundance analyses have been reported." + Comparison of abundances for the PAGB and RGB stars may reveal changes imposed by the evolution beyond the RGB: such changes are not necessarily attributable exclusively to internal nucleosynthesis and dredge-up., Comparison of abundances for the PAGB and RGB stars may reveal changes imposed by the evolution beyond the RGB; such changes are not necessarily attributable exclusively to internal nucleosynthesis and dredge-up. + It was in the spirit of comparing the compositions of the PAGB and RGB stars that we undertook our analysis., It was in the spirit of comparing the compositions of the PAGB and RGB stars that we undertook our analysis. + For the RGB stars. we use results kindly provided in advance of publication by Carretta (2008. private communication).," For the RGB stars, we use results kindly provided in advance of publication by Carretta (2008, private communication)." + 0.2 em PAGB stars because they are rapidly evolving are understandably rare in globular clusters., 0.2 cm PAGB stars because they are rapidly evolving are understandably rare in globular clusters. + At spectral types of F and G. a few luminous variables are known.," At spectral types of F and G, a few luminous variables are known." + These are sometimes referred to as Type II Cepheids., These are sometimes referred to as Type II Cepheids. + Abundance analyses of cluster variables have been reported. for example. for one or two stars in the clusters M2. M5. MIO. and M28 (Gonzalez Lambert 1997: Carney. Fry Gonzalez 1998).," Abundance analyses of cluster variables have been reported, for example, for one or two stars in the clusters M2, M5, M10, and M28 (Gonzalez Lambert 1997; Carney, Fry Gonzalez 1998)." + These stars have (DVJo of 0.5-0.6 rather than the (22Ευ=0.28 of the M79 discovery., These stars have $B-V)_0$ of 0.5-0.6 rather than the $(B-V)_0=0.28$ of the M79 discovery. +" At even earlier spectral types than A. the PAGB stars in globular stars are Widely referred to as ""UV-bright! stars (see review by Moehler 2001)."," At even earlier spectral types than A, the PAGB stars in globular stars are widely referred to as `UV-bright' stars (see review by Moehler 2001)." + Three such B-type stars have been subject to an abundance analvsis - see Thompson et al. (, Three such B-type stars have been subject to an abundance analysis - see Thompson et al. ( +2007).,2007). + 0.2 em In this paper. we present the abundance analysis of the M79 PAGB star and compare its composition to that of the cluster’s," 0.2 cm In this paper, we present the abundance analysis of the M79 PAGB star and compare its composition to that of the cluster's" +A disk powered by a black hole through magnetic connection shows interesting features different. [rom Chat of a standard accretion disk (Li2000b.2001).,"A disk powered by a black hole through magnetic connection shows interesting features different from that of a standard accretion disk \citep{li00b,li01}." +". The energy radiated by the disk comes from regions closer to the center of the disk. the radiation flix decreases more rapidly with radius. which approaches r?"" at large radii (compared to r* for a standard accretion disk)."," The energy radiated by the disk comes from regions closer to the center of the disk, the radiation flux decreases more rapidly with radius, which approaches $r^{-3.5}$ at large radii (compared to $r^{-3}$ for a standard accretion disk)." + This implies (hat (he radiation spectrum of a disk with magnetic connection observed by a distant observer will also be different [rom that of a standard accretion disk., This implies that the radiation spectrum of a disk with magnetic connection observed by a distant observer will also be different from that of a standard accretion disk. + Interestingly. the most recent observation of the nearby bright Sevlert 1 ealaxv AICG6-30-15 reveals an extremely. broad and red-shilted Fe ἵνα line indicating its origin [rom the very most central regions of the accretion disk (Wilmsetal.2001).," Interestingly, the most recent observation of the nearby bright Seyfert 1 galaxy MCG–6-30-15 reveals an extremely broad and red-shifted Fe $\alpha$ line indicating its origin from the very most central regions of the accretion disk \citep{wil01}." +. To explain the observed spectrum a very steep emissivity profile with index a=4.3—5.0 is required [a steep emissivitv lor the same galaxy. ancl Mrk-766 is also reported by (2001)]]., To explain the observed spectrum a very steep emissivity profile with index $\alpha = 4.3-5.0$ is required [a steep emissivity for the same galaxy and Mrk-766 is also reported by \citet{bra01}] ]. + Such a steep enissivity is very. difficult to be explained within the framework ol a standard accretion disk. but the magnetic connection between a black hole and a disk nav provide a natural explanation (Wilmsetal.2001).," Such a steep emissivity is very difficult to be explained within the framework of a standard accretion disk, but the magnetic connection between a black hole and a disk may provide a natural explanation \citep{wil01}." +. Though there is vet another possible explanation in term of the magnetic connection between the inner boundary of the disk ancl ihe plunging material inside the inner boundary (Ixrolik1999;Gammie1999:Agol 2000).. Che issue remains controversial (Paczviski2000) and the most recent numerical simulations show that (he stress produced by such a magnetic connection sensitively depends on the thickness of the disk: the stress in (he plunging region is signilicantlv reduced as the thickness of the disk decreases (Ilawley&INrolik2001a.b:Armitage.Revnold.2001:Hawley2001) important.," Though there is yet another possible explanation in term of the magnetic connection between the inner boundary of the disk and the plunging material inside the inner boundary \citep{kro99,gam99,ago00}, the issue remains controversial \citep{pac00} and the most recent numerical simulations show that the stress produced by such a magnetic connection sensitively depends on the thickness of the disk: the stress in the plunging region is significantly reduced as the thickness of the disk decreases \citep{haw01b,haw01c,arm01,haw01a} ." +.. In thisLeHer we use a simple model to demonstrate the observational signatures of ihe magnetic connection between a black hole and a disk., In this we use a simple model to demonstrate the observational signatures of the magnetic connection between a black hole and a disk. + We assume that an axisvimnmietric magnetic field connects à black hole to à non-accretion disk from the inner boundary of the disk to a circle with radius rj., We assume that an axisymmetric magnetic field connects a black hole to a non-accretion disk from the inner boundary of the disk to a circle with radius $r_b$. +" As in the ease of a standard accretion disk. the inner boundary ol the disk is assumed to be at the marginally stable orbit with radius r,,;."," As in the case of a standard accretion disk, the inner boundary of the disk is assumed to be at the marginally stable orbit with radius $r_{ms}$." + We will caleulate (he radiation flux. the emissivity index. aud the radiation spectrum observed by a distant observer for various distribution of the magnetic field between ως and ry on the disk.," We will calculate the radiation flux, the emissivity index, and the radiation spectrum observed by a distant observer for various distribution of the magnetic field between $r_{ms}$ and $r_b$ on the disk." + We will compare the results with that of a standard accretion disk and look for the robustness ol the observational signatures of the magnetic connection., We will compare the results with that of a standard accretion disk and look for the robustness of the observational signatures of the magnetic connection. +For the central density. appearing here. 0.. we use equation (7). evaluated at the cloud boundary: Finally. a itself max be written in terms of standard variables by. using equation (8): We now consiler a sequence of isothermal spheres of fixed mass. all embedded in the same external pressure.,"For the central density appearing here, $\delta_c$, we use equation (7), evaluated at the cloud boundary: Finally, $\alpha$ itself may be written in terms of standard variables by using equation (8): We now consider a sequence of isothermal spheres of fixed mass, all embedded in the same external pressure." + We may describe each structure using (he new. nondimensional variables.," We may describe each structure using the new, nondimensional variables." + The sequence is characterized bv a single parameter. the center-to-edge density contrast: we shall denote this ratio as 2.," The sequence is characterized by a single parameter, the center-to-edge density contrast; we shall denote this ratio as $\beta$." + From equation (1). 2 can alsobe written as since vc(£) is a known [unction. there is a one-to-one correspondence between our fundamental parameter 3 and £.. the old. nondimensional radius.," From equation (7), $\beta$ can alsobe written as Since $\psi (\xi)$ is a known function, there is a one-to-one correspondence between our fundamental parameter $\beta$ and $\xi_\circ$ , the old, nondimensional radius." + The potential 6 increases monotonically with £. so 53 likewise increases with £..," The potential $\psi$ increases monotonically with $\xi$, so $\beta$ likewise increases with $\xi_\circ$." + The lowest value of 3 is unitv. corresponding to0.," The lowest value of $\beta$ is unity, corresponding to." +. The internal velocity dispersion a varies along our sequence., The internal velocity dispersion $\alpha$ varies along our sequence. + We may (rack this change through equation (15)., We may track this change through equation (15). + Thus. for each selected 2. we first [ind @ from equation (16).," Thus, for each selected $\beta$, we first find $\psi_\circ$ from equation (16)." + From knowledge of the function c(£). we find the corresponding £.. as well as (ας.," From knowledge of the function $\psi (\xi)$, we find the corresponding $\xi_\circ$, as well as $\left(d\psi/d\xi\right)_\circ$." + Equation (15) then vields a., Equation (15) then yields $\alpha$. + It is equally streüghtForward to obtain the internal density. profile. 0(À). of any mocdel.," It is equally straightforward to obtain the internal density profile, $\delta (\lambda)$, of any model." + Knowing ος and a. equation (14) gives the central density. 9..," Knowing $\psi_\circ$ and $\alpha$ , equation (14) gives the central density, $\delta_c$." + Proceeding outward. equation (13) LOgives the value of hj€ correspondinge to each A.," Proceeding outward, equation (13) gives the value of $\xi$ corresponding to each $\lambda$." + Againe usinge Sc(£). equation (7) vields the density ratio.exp(c).," Again using $\psi (\xi)$, equation (7) yields the density ratio,." +. When we get to the edge.A... we find thatexp(—e.)..in agreement wilh equation (16).," When we get to the edge, we find that,in agreement with equation (16)." + Figure 1 displavs graphically (he change of the clouds structure as a function of 3., Figure 1 displays graphically the change of the cloud's structure as a function of $\beta$. + llere we have plotted the radius. A(2). of selected massshells.," Here we have plotted the radius, $\lambda (\beta)$ , of selected massshells." + As expected. a shell in the deep interior monotonically shrinks.," As expected, a shell in the deep interior monotonically shrinks." + Other shells. however. turn around.," Other shells, however, turn around." + With rising ». an, With rising$\beta$ an +Activation energy barriers are typically assumed to mediate key reactions in the methanol system.,Activation energy barriers are typically assumed to mediate key reactions in the methanol system. + However. these barriers were not considered in the study of Barzel&Biham(2007b).," However, these barriers were not considered in the study of \cite{barzel2}." +. In order to test the new modified-rate method against a system with activation energies. the methanol system of Stantchevaetal.(2002) is employed: see Tables | and 3.," In order to test the new modified-rate method against a system with activation energies, the methanol system of \cite{stant2} is employed; see Tables 1 and 3." + Stantchevaetal.(2002) used oa hybrid master-equation/rate-equation method to reproduce Monte Carlo results., \cite{stant2} used a hybrid master-equation/rate-equation method to reproduce Monte Carlo results. + The former approach assumed that the very reactive speciesH. O. OH. HCO and CHO behave stochastically. but that species with reactions requiring activation energies (CO and H2CO). which typically achieve large populations. could be treated using the standard rate equations.," The former approach assumed that the very reactive speciesH, O, OH, HCO and $_3$ O behave stochastically, but that species with reactions requiring activation energies (CO and $_2$ CO), which typically achieve large populations, could be treated using the standard rate equations." + The cut-offs for stochastic species were set to values |. 2. or 3.," The cut-offs for stochastic species were set to values 1, 2, or 3." +"4 Three basic scenarios were considered. characterised by the accretion fluxes of H. O. and CO atoms/molecules. mimicking conditions in interstellar clouds of “low”. ""intermediate"". and ""high"" density."," Three basic scenarios were considered, characterised by the accretion fluxes of H, O, and CO atoms/molecules, mimicking conditions in interstellar clouds of “low”, “intermediate”, and “high” density." + The models were run to 1000 yr of evolution. assuming 10° binding sites per grain and a grain temperature of 10 K. Importantly. Stantcheva et al.," The models were run to 1000 yr of evolution, assuming $10^6$ binding sites per grain and a grain temperature of 10 K. Importantly, Stantcheva et al." + allowed diffusion of hydrogen atoms around the grain surface to occur via quantum tunnelling between binding sites. rendering the rates of hydrogen reaction substantially faster than would be the case assuming only thermal hopping.," allowed diffusion of hydrogen atoms around the grain surface to occur via quantum tunnelling between binding sites, rendering the rates of hydrogen reaction substantially faster than would be the case assuming only thermal hopping." + Katzetal.(1999) suggested that such. quantum-tunnelling effects do not govern surface diffusio rates., \cite{katz} suggested that such quantum-tunnelling effects do not govern surface diffusion rates. + However. the purpose of the comparison is. in this case. to test the method against a reliable standard. rather than to provide an accurate reproduction of grain-surface chemistry m interstellar clouds.," However, the purpose of the comparison is, in this case, to test the method against a reliable standard, rather than to provide an accurate reproduction of grain-surface chemistry in interstellar clouds." + More rigorous application of the new modification methods to interstellar cloud conditions will follow in future., More rigorous application of the new modification methods to interstellar cloud conditions will follow in future. + Of the two methods utilised by Stantcheva et al..," Of the two methods utilised by Stantcheva et al.," + the Monte Carlo results should be considered the most reliable. although the differences are typically not large.," the Monte Carlo results should be considered the most reliable, although the differences are typically not large." + However. the Monte Carlo simulations use only integer values. so it is not possible to obtain accurate estimates of average population sizes when those values are close to unity.," However, the Monte Carlo simulations use only integer values, so it is not possible to obtain accurate estimates of average population sizes when those values are close to unity." + Comparison with species of very low abundance relies on the results of the master-equation/rate-equation hybrid method alone., Comparison with species of very low abundance relies on the results of the master-equation/rate-equation hybrid method alone. +" For convenience. this method will from now on be refered to simply as the ""master equation? method. with the understanding that it is in fact a hybrid scheme."," For convenience, this method will from now on be refered to simply as the “master equation” method, with the understanding that it is in fact a hybrid scheme." + Table 4 shows two sets of results from Stantcheva et al.:, Table 4 shows two sets of results from Stantcheva et al.; +" the column headed ""Master Equation 22211"" indicates results from their model using cut-offs for stochastic species H. O. OH. HCO. CHO of 2. 2. 2. |. I. respectively."," the column headed “Master Equation 22211” indicates results from their model using cut-offs for stochastic species H, O, OH, HCO, $_3$ O of 2, 2, 2, 1, 1, respectively." + Columns in tables 5 and 6 are labelled similarly., Columns in tables 5 and 6 are labelled similarly. + Stantcheva et al., Stantcheva et al. + also ran models with lower cut-offs: the results of the models with the highest cut-offs are quoted. on the assumption that these are the most accurate.," also ran models with lower cut-offs; the results of the models with the highest cut-offs are quoted, on the assumption that these are the most accurate." + Tables 4 — 6 also show rate-equation solutions for each regime., Tables 4 – 6 also show rate-equation solutions for each regime. + Whilst a few rate-equation results are within an order of magnitude of the exact value. many populations vary wildly from the Monte Carlo and master-equation results. including important species such as formaldehyde. methanol and CO>.," Whilst a few rate-equation results are within an order of magnitude of the exact value, many populations vary wildly from the Monte Carlo and master-equation results, including important species such as formaldehyde, methanol and $_2$." + In fact. the rate-equation results shown in Tables 4 — 6 differ from the values quoted by Stantchevaetal.(2002)... which actually correspond to activation energies of 2500 K. for the H + CO and H + H:CO reactions. rather than the stated 2000 K. Nevertheless. 1t is clear that the rate equations produce à very poor match to the exact solutions of each system.," In fact, the rate-equation results shown in Tables 4 – 6 differ from the values quoted by \cite{stant2}, which actually correspond to activation energies of 2500 K, for the H + CO and H + $_2$ CO reactions, rather than the stated 2000 K. Nevertheless, it is clear that the rate equations produce a very poor match to the exact solutions of each system." +" In Table 4. the column headed ""Method A"" shows low-density results using the simple competition scheme of Section 5."," In Table 4, the column headed “Method A” shows low-density results using the simple competition scheme of Section 5." + Aside from O*. CO. H»CO and COs. all species show an excellent match to the Monte Carlo and/or master equation results.," Aside from $_2$, CO, $_2$ CO and $_2$, all species show an excellent match to the Monte Carlo and/or master equation results." + The worst. CO». is only a factor of ~ 1.4 greater than the master-equation value; a great improvement over the rate-equations.," The worst, $_2$, is only a factor of $\sim$ 1.4 greater than the master-equation value; a great improvement over the rate-equations." + Table 5 shows results for the intermediate density case., Table 5 shows results for the intermediate density case. + Here. the simple competition scheme produces a good match with all species except CO». which ts inaccurate by a factor similar to that of the low-density case.," Here, the simple competition scheme produces a good match with all species except $_2$, which is inaccurate by a factor similar to that of the low-density case." + The quality of the match for all other species is not in general as good as for the case. but is still quite acceptable.," The quality of the match for all other species is not in general as good as for the low-density case, but is still quite acceptable." + Table 6 shows results for the high density case., Table 6 shows results for the high density case. + The results of the simple competition scheme are generally within an order of magnitude of the Monte Carlo and/or master-equation results. and some are rather closer: O» is very accurate.," The results of the simple competition scheme are generally within an order of magnitude of the Monte Carlo and/or master-equation results, and some are rather closer; $_2$ is very accurate." + In all regimes. method A provides at least a reasonable match to the exact results. and great improvements over the rate-equation method.," In all regimes, method A provides at least a reasonable match to the exact results, and great improvements over the rate-equation method." + However. there are still discrepancies that cannot be explained merely as computational inaccuracies between the different methods.," However, there are still discrepancies that cannot be explained merely as computational inaccuracies between the different methods." + Further refinements to the treatment of surface processes are required to match the exact results., Further refinements to the treatment of surface processes are required to match the exact results. + Section 5 outlined a simple competition scheme that introduced an efficiency to the surface production rates., Section 5 outlined a simple competition scheme that introduced an efficiency to the surface production rates. + The scheme assumes that if one or other particle evaporates before reaction can occur. then the production process is over.," The scheme assumes that if one or other particle evaporates before reaction can occur, then the production process is over." + However. this is an over-simplification; the final rate must take into account all opportunities for an individual. waiting particle to react before it is otherwise removed.," However, this is an over-simplification; the final rate must take into account all opportunities for an individual, waiting particle to react before it is otherwise removed." + Consider particle A waiting on a gram for its reaction partner. B.," Consider particle $A$ waiting on a grain for its reaction partner, $B$." + Particle B accretes onto the grain. and if A and B meet before evaporation of either species occurs. then they react.," Particle $B$ accretes onto the grain, and if $A$ and $B$ meet before evaporation of either species occurs, then they react." +" If particle A itself evaporates before reaction can occur. then the process is over. as the part of the production rate associated with this scenario. Rj,;,,0D)ΝΑ). presumes the presence of particle A on the grain."," If particle $A$ itself evaporates before reaction can occur, then the process is over, as the part of the production rate associated with this scenario, $R_{form}(B) \cdot \langle N(A) \rangle$, presumes the presence of particle $A$ on the grain." + But if particle B evaporates. particle A is still present and waiting for another particle B to acerete. of which there is asteady stream.," But if particle $B$ evaporates, particle $A$ is still present and waiting for another particle $B$ to accrete, of which there is asteady stream." + If particle A can remain on the gram for another accretion timescale of B. without evaporating. then the reaction competition process may begin again.," If particle $A$ can remain on the grain for another accretion timescale of $B$ , without evaporating, then the reaction competition process may begin again." + Two competition processes must therefore be considered: competition between reaction and evaporation ofA or B; and competition between accretion of B and evaporation of A., Two competition processes must therefore be considered: competition between reaction and evaporation of $A$ or $B$ ; and competition between accretion of $B$ and evaporation of $A$ . + lgnoring interference from other accreting particles.," Ignoring interference from other accreting particles," +diffusiou-loss equation for sonie assuned system. mass and diffusion loss coefficient. but in either the thiu- or thick-tareet case the required timescale is only of order vears. and over the lifetiue of the black hole erowth. a substantial total amount of naterial could be processed throteh the cosmic rav region: at anv one iusant of tiue we only need to be observing a small mass of spallation-cuhanced eas which could lave teο substantially cuhauced in its Cr abundance.,"diffusion-loss equation for some assumed system mass and diffusion loss coefficient, but in either the thin- or thick-target case the required timescale is only of order years, and over the lifetime of the black hole growth, a substantial total amount of material could be processed through the cosmic ray region: at any one instant of time we only need to be observing a small mass of spallation-enhanced gas which could have been substantially enhanced in its Cr abundance." + T119 scenario ds al effective wav of explüniug the apparent amudance eniucements that we see without requiring extreue conditions of cosmic rav production., This scenario is an effective way of explaining the apparent abundance enhancements that we see without requiring extreme conditions of cosmic ray production. +" Tuteractions of cosnüc ravs with the accreting gas are expected to produce secondary 5-ràw cluission (7) aud ?7.equations11-12). estimated the x-ray flux above LOO MeV. for Sevtert galaxies if spallation is ουήπιο,", Interactions of cosmic rays with the accreting gas are expected to produce secondary $\gamma$ -ray emission \citep{dermer86a} and \citet[][equations 11-12]{skibo97} estimated the $\gamma$ -ray flux above 100 MeV for Seyfert galaxies if spallation is occurring. + We have recaleulated the expected οταν fux specifically for 11051 following ?).., We have recalculated the expected $\gamma$ -ray flux specifically for 4051 following \citet{skibo97}. + The estimate considers only production of *s-ravs following neutral pion creation. as calculated for the AGilky Way dy 7). whose code we use for the calculation. modified to incorporate LL of ?)..," The estimate considers only production of $\gamma$ -rays following neutral pion creation, as calculated for the Milky Way by \citet{dermer86a}, whose code we use for the calculation, modified to incorporate 4 of \citet{skibo97}." + The calculation asstunes tje “thick. arect™ case where all proteis are absorbed. audxd hence this provides a maxilui 2-rav bhuuinositv or a eiven output ACN proton lhuninositv.," The calculation assumes the “thick target” case where all protons are absorbed, and hence this provides a maximum $\gamma$ -ray luminosity for a given output AGN proton luminosity." +" ÁAssunüug that Lo,=Lp. the predicted 5-rav flux is ΕςLOOMeV)~δν10""(071z)photonssEcui7 nfor protou spectral iudex Q=2.| and proton low-energv cut-off MeV. About ppercent of the input cosmic rav eneregev is radiated as 5-ravs."," Assuming that $L_{\rm\scriptscriptstyle CR} = L_{bol}$, the predicted $\gamma$ -ray flux is $F (> 100\, {\rm MeV}) \simeq 3.3 \times 10^{-8} (\Omega/4\pi){\rm photons\, s^{-1}\, cm^{-2}\,}$ for proton spectral index $\alpha=2.4$ and proton low-energy cut-off MeV. About percent of the input cosmic ray energy is radiated as $\gamma$ -rays." + For this prooll input spectrum. the 5-rav spectrum ds hareom with an approximately power-lav form cwith photon index 2.1 in the range LOOMAIeV<2.LOO CCGeV. 7?) discuss the detection limit of 1CFern’ Large Area Telescope aud for ο~0.9 the predicted flux of 11051. falls just beQW the 10e deection Tit frou he 3anouth skv survey data ~3.5.«10ESN“photousstem2? for a ⋅ ⋅↴∖↴≺∏∐⋅∩∖∪↕↑↕∐↴∖↴↴∖↴⋉∖↸⊳⊓⋅⋯⊔⋜∐∐↧≼∣⋜↧↕⋯⊳⊓↸⊳↸⊳∪∪↥⋅≼∐∐⋜↧ À(?7)..," For this proton input spectrum, the $\gamma$ -ray spectrum is harder, with an approximately power-law form with photon index 2.1 in the range $ < E_\gamma < 100$ GeV. \citet{abdo09a} discuss the detection limit of the Large Area Telescope and for $\Omega \sim 0.9$ the predicted flux of 4051 falls just below the $\sigma$ detection limit from the 3-month sky survey data $\sim 3.5 \times 10^{-8} {\rm photons\, s^{-1}\, cm^{-2}\,} $ for a source of this spectrum and Galactic coordinates \citep{abdo09a}." +"tes ""thi the ACN.ΠΙΟ tareet in which protons esca20 ower global coveri18o factors. lower COSLUG ray huuinosities achieved through longer spallation timescales or the possiülitv of internitent cosuuic ταν output all weaken the possible coustyants available from imeasurenienuts of 2-rav flux."," A “thinner” target in which protons escape the AGN, lower global covering factors, lower cosmic ray luminosities achieved through longer spallation timescales or the possibility of intermittent cosmic ray output all weaken the possible constraints available from measurements of $\gamma$ -ray flux." + Radio svuchrotron eLulsson nuelt also be expected. partly as a result of secondary electron creation but more significantly through primary cosnüc ray electrons soirerated as part of the process hat accelerates tle protons;," Radio synchrotron emission might also be expected, partly as a result of secondary electron creation but more significantly through primary cosmic ray electrons generated as part of the process that accelerates the protons." + Auv estimate of the svuchrotrou emission is extremevo uncertaiu. as we dont kuow the efhiciency of electron acceleraion compared with proton acceleration. we dont know the maeietic field sreneth axd the elecτος thoniselves are expectec to suffer sieuificait donization losses propagating through the material as well as adiabatie ane radiative losses through brenisstrahluug. svuchrotrou axd inverse Compton scatterjus.," Any estimate of the synchrotron emission is extremely uncertain, as we don't know the efficiency of electron acceleration compared with proton acceleration, we don't know the magnetic field strength and the electrons themselves are expected to suffer significant ionization losses propagating through the material as well as adiabatic and radiative losses through bremsstrahlung, synchrotron and inverse Compton scattering." + Furthermore. aux conrpact svuchrotron-chutting reegiou would be seltabsorbed. siguificautlv reducing the «cetect-ability of anv svuchrotro1 enisson.," Furthermore, any compact synchrotron-emitting region would be self-absorbed, significantly reducing the detect-ability of any synchrotron emission." + We oO nof here solve t1e diffusion-loss equation for electrous. but we can already pace a strimecut nuit frou siuple estinates of the radiative losses.," We do not here solve the diffusion-loss equation for electrons, but we can already place a stringent limit from simple estimates of the radiative losses." + The ratio of svuchrotrou o inverse Compton losses is given approxinatelv |w the ratio of energv deusities in magnetic fiokls. muguad. radiation. a.. so if these were the dominant energv loss niechanisnis. we would expect an inteerated svuchrotrou Iuniuositv ucelecting svuchrotrou selt-absorptiou for the moment.," The ratio of synchrotron to inverse Compton losses is given approximately by the ratio of energy densities in magnetic fields, $u_{mag}$, and radiation, $u_\gamma$, so if these were the dominant energy loss mechanisms, we would expect an integrated synchrotron luminosity neglecting synchrotron self-absorption for the moment." + If other mechauisius lead to comparable enerev losses. the svuchrotron bhuuimositv would be lower.," If other mechanisms lead to comparable energy losses, the synchrotron luminosity would be lower." + 11051 Lhas an stecp spectrum radio core with flux densi v2anuuJyw at (ιν when measured at ECSOution l.laaresec (7).. corresponding to a spatia resolution ype.," 4051 has an steep spectrum radio core with flux density mJy at GHz when measured at resolution arcsec \citep{ho01a}, corresponding to a spatial resolution pc." + The core has been claimed to be mmareiually detected at EVN resolution (?).., The core has been claimed to be marginally detected at EVN resolution \citep{giroletti09a}. +" Tre core radio hhnunositv ofL1051 iuteerate Loup to 100CCGIIz is cotxoxinatelv 3«4107 ress +. so we can place approxinate upper iut on the magnetic fied streneth if we require L4, to be ess thaw this value."," The core radio luminosity of4051 integrated up to GHz is approximately $3 \times 10^{37}$ $^{-1}$ , so we can place an approximate upper limit on the magnetic field strength if we require $L_{\rm sync}$ to be less than this value." + For LcLigc1 !| and an euergv denusitv iu radiation at p= ceni of àzTOO 5m we find duas<0.002 3 COYTORSDouding to ainagnetic flux density D«0.2 Gauss.," For $L_{\rm electron} \simeq L_{\rm bol} \simeq 10^{43}$ $^{-1}$ and an energy density in radiation at $r=2 \times 10^{14}$ cm of $u \simeq 700$ $^{-3}$ we find $u_{\rm mag} < 0.002$ $^{-3}$ , corresponding to amagnetic flux density $B < 0.2$ Gauss." + Iu the svuchrotrou- regime a discrepant radio flix could oulv be produced if D were hieher than this, In the synchrotron-optically-thin regime a discrepant radio flux could only be produced if $B$ were higher than this +Furally. all of our conclusloli regarding he effectiveness of stellar wind to1QUes relies on he particular formulation for the torque from Matt&Dudritz (2008a}.,"Finally, all of our conclusions regarding the effectiveness of stellar wind torques relies on the particular formulation for the torque from \citet{mattpudritz08II}." +. While this is fi0 lnost appropriate forunlation for low-nass-star winds that exists iu he literaure. there remain a few ope1 questions about its use 1u the preseut work.," While this is the most appropriate formulation for low-mass-star winds that exists in the literature, there remain a few open questions about its use in the present work." + First of all. that formulation Was derived for au isolated star. and we have simply adclec this torque to our model. wch also assumes the preseice of an accretion disk aud associated Interaction torques in a non-«self-consisent wav.," First of all, that formulation was derived for an isolated star, and we have simply added this torque to our model, which also assumes the presence of an accretion disk and associated interaction torques in a non-self-consistent way." + It is not vet clear how t1e presence of a disk will ifence the stellar wixl. aud i particular how it may ithence the depenudeice of the stellar wind torque οji parameters (e.g. the vaue of in ).," It is not yet clear how the presence of a disk will influence the stellar wind, and in particular how it may influence the dependence of the stellar wind torque on parameters (e.g., the value of $m$ )." + Also. the formulation of M;tt&Pudritz(20082) was based upou sinulatiois of a star with a sine)lar spin rate (f 0.1) and a sobu-like wind acceleration mechanisin.," Also, the formulation of \citet{mattpudritz08II} was based upon simulations of a star with a singular spin rate $f=0.1$ ) and a solar-like wind acceleration mechanism." +" It is eucouragiug that a study of auglar momentum flow frou, massive stars (Ud-Doulactal.2008. 2009).. which iuchludes a very different wind acceleration mechanisin and a range of spin rates. fotnd a smnilar power law relationship for the torque. but more studies are warranted."," It is encouraging that a study of angular momentum flow from massive stars \citep{uddoula3ea08, uddoula3ea09}, which includes a very different wind acceleration mechanism and a range of spin rates, found a similar power law relationship for the torque, but more studies are warranted." + Lastly. all stellar wind torque formmlations in the literature are based upon simple (e.8.. dipolar) magnetic ecometrics. while we know that T Tauri stars possess complex magnetic field structures.," Lastly, all stellar wind torque formulations in the literature are based upon simple (e.g., dipolar) magnetic geometries, while we know that T Tauri stars possess complex magnetic field structures." + Iu future work. it will be important to determine whether the presence of au accretion disk. more complex mageic ecometrics. or different wind acceleration mechajisnis will senifüeaulty cuhance or suppress the stellar wind orque relative to the formulation of Matt&Pudritz(2008a) adopted here.," In future work, it will be important to determine whether the presence of an accretion disk, more complex magnetic geometries, or different wind acceleration mechanisms will significanlty enhance or suppress the stellar wind torque relative to the formulation of \citet{mattpudritz08II} adopted here." + Mc1 work rendus to develop the necessary theory. to nuprove t16 precision and uuuber of observational luüeasurenieifs of relevant system parameters. and ultinacly to nuderstaud the observed distributions aud evoutkn of κcllar spin rates. but the idea that powerful «telar winds nay extract significant angular momentum from accretingc» stars renuadus a pronisiuec» scenario.," Much work remains to develop the necessary theory, to improve the precision and number of observational measurements of relevant system parameters, and ultimately to understand the observed distributions and evolution of stellar spin rates, but the idea that powerful stellar winds may extract significant angular momentum from accreting stars remains a promising scenario." + We thauk the auonviuous referee for useful remarks on tha ognuuiuscripte, We thank the anonymous referee for useful remarks on the manuscript. +" SPAL was supported bv an appointment to the NASA Postdoctoral Program at Alues Researei Center, aduuinistered by Oak Ridge Associated Universities through a contract with NASA. aud by the ERC through eraut 207130 STARS2 (http:/wwwestars2.cu)."," SPM was supported by an appointment to the NASA Postdoctoral Program at Ames Research Center, administered by Oak Ridge Associated Universities through a contract with NASA, and by the ERC through grant 207430 STARS2 (http://www.stars2.eu)." + TPG acknowledges support from NzASA’s Origins of Solar Svsteiis program via WBS δ110τὸ.02.07.0189., TPG acknowledges support from NASA's Origins of Solar Systems program via WBS 811073.02.07.01.89. + (Bxz10!G) , $B\lesssim 10^{10}\ {\rm G}$ +for the fact that inner and outer disk abundances refer to different elements. and (i1) the age-metallicity relation. which reconciles abundances in young HII regions with metallicities of the stars on the red giant branch which we have assumed to be at least 8 Gyr old.,"for the fact that inner and outer disk abundances refer to different elements, and (ii) the age-metallicity relation, which reconciles abundances in young HII regions with metallicities of the stars on the red giant branch which we have assumed to be at least 8 Gyr old." + Despite numerous results pointing to a flat age-metallicity relationship in the Solar neighborhood (e.g.?).. these works find that the relation steepens for the oldest stars.," Despite numerous results pointing to a flat age-metallicity relationship in the Solar neighborhood \citep[e.g.][]{edvardsson93}, these works find that the relation steepens for the oldest stars." + Assuming that the difference in metallicities of young HII regions and 8—12 Gyr old RGBstars Is ~0.3—0.5 dex and |o/Fe]20.0—0.3. it is possible to account for the abundance difference of 0.5—0.8 dex in the transition region in Figure 7...," Assuming that the difference in metallicities of young HII regions and $8-12$ Gyr old RGBstars is $\sim0.3-0.5$ dex and $\alpha$ $=0.0-0.3$, it is possible to account for the abundance difference of $0.5-0.8$ dex in the transition region in Figure \ref{gradient}." + At the high end O.Sdex). thismakeshighinnerdiskmetallicitiesconsistentwiththeverylotvilihedoutarsdwdteatv," At the high end $0.8$ dex), this makes high inner disk metallicities consistent with the very low abundances we derive." + Furthermore. in deriving metallicities of regions. ? and ? use strong-line abundance indicators.," Furthermore, in deriving metallicities of regions, \citet{edmundspagel84} and \citet{webstersmith83} use strong-line abundance indicators." + These have been shown to result in metallicities which are a factor of two (0.3 dex) higher than nebular (7; based) or stellar abundances and likely enhance the difference between the inner line) and outer disk (stellar) abundances in Figure 7.., These have been shown to result in metallicities which are a factor of two $0.3$ dex) higher than nebular $T_e$ based) or stellar abundances \citep{bresolin09b} and likely enhance the difference between the inner (strong-line) and outer disk (stellar) abundances in Figure \ref{gradient}. . + Similar abundance behavior (although of a smaller magnitude) is observed in the Milky Way., Similar abundance behavior (although of a smaller magnitude) is observed in the Milky Way. + In Figure 8 we reproduce Figure 4 from ?.., In Figure \ref{mw} we reproduce Figure 4 from \citet{carney05}. + Old open clusters are shown as diamonds (from??) and squares (from?)..," Old open clusters are shown as diamonds \citep[from][]{twarog97,chen03} and squares \citep[from][]{yong05}." + Cireles are three field red giants from ?.., Circles are three field red giants from \citet{carney05}. . + Cepheids from the works of ?.. ?.. 2..2.. and? are shown as crosses.," Cepheids from the works of \citet{andrievsky02a}, \citet{andrievsky02b}, \citet{andrievsky02c}, \citet{luck03}, and \citet{andrievsky04} are shown as crosses." + Cepheids follow the well-defined negative gradient out to ~15 kpe., Cepheids follow the well-defined negative gradient out to $\sim15$ kpc. + On the contrary. old stars exhibit steeper gradient in the inner disk (out to 10—12 kpc) and a flattening in the outermost parts.," On the contrary, old stars exhibit steeper gradient in the inner disk (out to $10-12$ kpc) and a flattening in the outermost parts." + In the region 10—15 kpe there is a clear disconnect of up to 0.4 dex between the Cepheid abundances and those of old stars. similar to what we observe in NGC 7793.," In the region $10-15$ kpc there is a clear disconnect of up to $0.4$ dex between the Cepheid abundances and those of old stars, similar to what we observe in NGC 7793." + We can summarize the results presented thus far as follows: 7793 ts primarily old with red giant branch stars as the NGCdominant stellar population and a small contribution from asymptotic red giant branch stars. (, We can summarize the results presented thus far as follows: (i) The outer disk of NGC 7793 is primarily old with red giant branch stars as the dominant stellar population and a small contribution from asymptotic red giant branch stars. ( +11) After the contamination from faint background galaxies has been taken into account. number counts of RGB stars. as well as those derived using a matched catalog of all stars in the field. extend out to ~I0— L. or ~10.5—11.5 kpe. (,"ii) After the contamination from faint background galaxies has been taken into account, number counts of RGB stars, as well as those derived using a matched catalog of all stars in the field, extend out to $\sim10-11'$ , or $\sim10.5-11.5$ kpc. (" +iit) The effective surface brightness profile derived from star counts traces the disk of NGC 7793 out to 9 disk scale lengths and is ~3 mag arcsec deeper than the surface photometry data of ?..,iii) The effective surface brightness profile derived from star counts traces the disk of NGC 7793 out to $9$ disk scale lengths and is $\sim3$ mag $^{-2}$ deeper than the surface photometry data of \citet{carignan85}. . + Any potential break in the light profilemay be associated with a specific stellar population. but we seeno evidence of a truncation in old stars. (," Any potential break in the light profilemay be associated with a specific stellar population, but we seeno evidence of a truncation in old stars. (" +iv) The metallicity gradient inthe outer disk of NGC 7793,iv) The metallicity gradient inthe outer disk of NGC 7793 +format in the 1999 Brugual-Charlot library for isochrone svuthesis spectral evolution (Charlot Druzual 1991: Druzual Charlot 1993). asstuuine solar abundance.,"format in the 1999 Bruzual-Charlot library for isochrone synthesis spectral evolution (Charlot Bruzual 1991: Bruzual Charlot 1993), assuming solar abundance." + The total B-banud Iuuiuositv of the dropout stellar population is whereρω) is the single burst D-baud lunimositv per uni solar mass.," The total B-band luminosity of the dropout stellar population is where$\ell_{B,do}(t)$ is the single burst B-band luminosity per unit solar mass." +" The Dhnuuinositv of the backeround old single burst population is Lpalf)μινεου). where f, Is a cocfiicicut of order unity that must be adjusted for aerecment with the observed B-baud Iuinosity of NGC 1172 (see below)."," The luminosity of the background old single burst population is $L_{B,old}(t) = f_m M_{*t} \ell_{B,old}(t)$ where $f_m$ is a coefficient of order unity that must be adjusted for agreement with the observed B-band luminosity of NGC 4472 (see below)." + The evolution of B-bancd bIuuiuosities aud D- colors are illustrated in Figure lb aud le for each stellar population aud for thei combined radiation., The evolution of B-band luminosities and B-V colors are illustrated in Figure 1b and 1c for each stellar population and for their combined radiation. +" For reference. the redslitt 5:=1 isshown attime t=6.19 Gyre (assunüug Z7,=65 laus |, O,,=0.3 and Ον= 0.7)."," For reference, the redshift $z = 1$ isshown attime $t = 6.19$ Gyrs (assuming $H_o = 65$ km $^{-1}$, $\Omega_m = 0.3$ and $\Omega_{\Lambda} = 0.7$ )." + We have chosen fin=1.35 so that the total B-haucd huuiuositv Lnes=Lnoa|LnasoTSs10 Lp. appropriate for NGC 1172 at distance d=17 Alpe.," We have chosen $f_m = 1.35$ so that the total B-band luminosity $L_{B,tot} = L_{B,old} + L_{B,do} = 7.89 \times 10^{10}$ $L_{B,\odot}$, appropriate for NGC 4472 at distance $d = 17$ Mpc." + Although the dropout population currently contributes about 15 percent of the total B-baud light (Lyἐν)1.2«.lol’! Lp.) its fractional contribution to the ealactic light is quite coustaut for redshifts 5:x1.," Although the dropout population currently contributes about 15 percent of the total B-band light $L_{B,do}(t_n) = 1.2 \times 10^{10}$ $L_{B,\odot}$ ), its fractional contribution to the galactic light is quite constant for redshifts $z \lta 1$." + The combined dololl population is onlv slightly bluer (bx (B.V)~ 0.03) than the old population and this difference is esseutially coustaut for :X1., The combined do+old population is only slightly bluer (by $\delta(B-V) \sim 0.03$ ) than the old population and this difference is essentially constant for $z \lta 1$. + Formally. the 2V of the combined population indicates an age ~8.5 Cors that is less than that of the old population alone. 12 Civis. but metallicity variatious could produce a similar color variation.," Formally, the $B - V$ of the combined population indicates an age $\sim 8.5$ Gyrs that is less than that of the old population alone, 12 Gyrs, but metallicity variations could produce a similar color variation." +" To estimate the equivalent width of IL? we asstuue that the line width is similar for both populations: where Ποιο f,CR) is the ratio of light. from the old to dropout population within projected radius RR2 normalized to the otal ratio of old to dropout light: eenerallv we :wsnne fy1l. correspoudiug to viewing the total liebt from both populations."," To estimate the equivalent width of $\beta$ we assume that the line width is similar for both populations: where Here $f_g(R)$ is the ratio of light from the old to dropout population within projected radius $R$ normalized to the total ratio of old to dropout light; generally we assume $f_g = 1$, corresponding to viewing the total light from both populations." +" As an illustratioj we use suele burst ο(1) from the 1999 Bruzual-Charlot (BC99) tables appropriate to the IMF of cach population,"," As an illustration, we use single burst $ew_{\beta}(t)$ from the 1999 Bruzual-Charlot (BC99) tables appropriate to the IMF of each population." +" As shown in Figme 2 he dropout population (with fo=1) reduces he apparent age of the old )pulation by ~5 Crs. be. ντε)mEMalt,5Cars). in agreement with observations hat correlate bluer colors with stronger I> (Forbes Pouman 1999)."," As shown in Figure 2 the dropout population (with $f_g = 1$ ) reduces the apparent age of the old population by $\sim 5$ Gyrs, i.e. $EW_{\beta,tot}(t_n) +\approx EW_{\beta,old}(t_n - 5~{\rm Gyrs})$ , in agreement with observations that correlate bluer colors with stronger $\beta$ (Forbes Ponman 1999)." +" For smaller fy=L/L. corresponding to viewing NGC 1172 within r,. he apparent age is reduced by ~8&5 Cars."," For smaller $f_g = 1/4$ , corresponding to viewing NGC 4472 within $r_e$, the apparent age is reduced by $\sim 8.5$ Gyrs." + Actual observations of galactic cores view a raction of both populations. ie.the appareut IL/ age is aperture-dependent.," Actual observations of galactic cores view a fraction of both populations, i.e.the apparent $\beta$ age is aperture-dependent." +" For my,22. etyuqu0) aud EW¢or(ty} ave iuscusitive to ny. Le. etyalt)=οστ) can be assmnaed."," For $m_u \gta 2$, $ew_{\beta,do}(t)$ and $EW_{\beta,tot}(t_n)$ are insensitive to $m_u$, i.e. $ew_{\beta,do}(t) = +ew_{\beta,old}(t)$ can be assumed." +" We also used the population code available at the Worthev website to determine both ΤΕνετ) aud (5.Vf) for a Salpeter dropout IME from m,=0.2 toa, =LO AL..."," We also used the population code available at the Worthey website to determine both $EW_{\beta,tot}(t_n)$ and $(B - V)_{tot}(t_n)$ for a Salpeter dropout IMF from $m_{\ell} = 0.2$ to $m_u = 10$ $M_{\odot}$." + For this model the IL/aud (B.V) ages are 8.5 and 9.5 Cor, For this model the $\beta$ and $(B - V)$ ages are 8.5 and 9.5 Gyrs with $f_g = 1$. +rs with procedures (Charlot. Worthev Bressan 1996: Worthev 1996).," These age uncertainties are consistent with $\sim 35$ ) errors inherent to population synthesis procedures (Charlot, Worthey Bressan 1996; Worthey 1996)." + Clearly. however. the coutributiou of cooling dropout stars to the spectra of elliptical galaxies can explain the relatively vouus ages inferred from ΤΠ.) observed iu some ellipticals even if the underline stars are very old.," Clearly, however, the contribution of cooling dropout stars to the spectra of elliptical galaxies can explain the relatively young ages inferred from $\beta$ observed in some ellipticals even if the underlying stars are very old." + Allimassive ellipticals contain cooling iuterstellar σας SO Lo coniparisous with eas-free galaxies can be made., All massive ellipticals contain cooling interstellar gas so no comparisons with gas-free galaxies can be made. + However. since the cooling flow mass dropout is centrally concentrated within +... the equivalent width of Ilo. Εν should increase toward galactic centers. in agreement with the observations of Conzalez (1995).," However, since the cooling flow mass dropout is centrally concentrated within $r_e$ , the equivalent width of $\beta$ , $EW_{\beta}$ , should increase toward galactic centers, in agreement with the observations of Gonzalez (1993)." + Dropout star formation should be accompanied by an euseuble of additional observations atsmall galactic radii: ILJ trou cooling clouds. enhanced X- surface brielitucss due to deuse. locally cooling regions. and lower appareut X-ray temperatures," Dropout star formation should be accompanied by an ensemble of additional observations atsmall galactic radii: $\beta$ from cooling clouds, enhanced X-ray surface brightness due to dense, locally cooling regions, and lower apparent X-ray temperatures" +"For the standard model of LTO. with the NOs star-forluaion law, BXX7and no winds. we find: Fiuallv. for the model of LTQ with tlie IK98 star-forination law. Dxp. and wiuds in starbursts. we fiud:","For the standard model of LTQ, with the K98 star-formation law, $B \propto \Sigma_g^{0.7}$, and no winds, we find: Finally, for the model of LTQ with the K98 star-formation law, $B \propto \rho^{0.5}$, and winds in starbursts, we find:" +" ~~110 (£z109 F’, log(fT0). TU Metzger et al.", $\sim1-10$ $t\ga 10^6$ $F_x$ $(t-T0)$ $T0$ Metzger et al. + 200δα. 2008b. Cannizzo Gehrels 2009. hereafter CGO9).," 2008a, 2008b, Cannizzo Gehrels 2009, hereafter CG09)." + For long GRBs (IGRBs). the early. steep rate of decay may be giving us information about the radial density distribution within the progenitor core (Kumar et al.," For long GRBs (lGRBs), the early, steep rate of decay may be giving us information about the radial density distribution within the progenitor core (Kumar et al." + 2008b). whereas the later decay may be governed by the outward expansion of the transient disk formed from the remnants of the progenitor (CGO9).," 2008b), whereas the later decay may be governed by the outward expansion of the transient disk formed from the remnants of the progenitor (CG09)." + As regards short GRBs (sGRBs). even if a small amount of material (~10°?10ο 1M.) is expelled during the NS-NS merger and later accreted in a disk. that would be sufficient to power a bright afterglow. which may also be strongly influenced by the effects of r-process nucleosynthetie heating in the neutron rich material that becomes the disk (Metzger et al.," As regards short GRBs (sGRBs), even if a small amount of material $\sim 10^{-5} - 10^{-4}\msun$ ) is expelled during the NS-NS merger and later accreted in a disk, that would be sufficient to power a bright afterglow, which may also be strongly influenced by the effects of r-process nucleosynthetic heating in the neutron rich material that becomes the disk (Metzger et al." + 2010)., 2010). + Zhang et al. (, Zhang et al. ( +2006) present a schematic for the decaying GRB light curve as seen by the XRT onSwift.,2006) present a schematic for the decaying GRB light curve as seen by the XRT on. + The decay is traditionally shown in logF—logf., The decay is traditionally shown in $\log F - \log t$. +" There are four basic power-law decay (F«Xf£ "") regimes: (1) a steep decline following the prompt emission with oj~3 out to 107.—LO? s. (iD) a plateau with ay~0.5 out to 107—104 s. (ii à steepening with ay;c1.2 out to 10—10? s. and (iv) à further steepening at late times (not always seen) with ayy&2 "," There are four basic power-law decay $F \propto t^{-\alpha}$ ) regimes: (i) a steep decline following the prompt emission with $\alpha_{\rm I} \simeq 3$ out to $10^2-10^3$ s, (ii) a plateau with $\alpha_{\rm II} \simeq 0.5$ out to $10^3-10^4$ s, (iii) a steepening with $\alpha_{\rm III} \simeq 1.2$ out to $10^4-10^5$ s, and (iv) a further steepening at late times (not always seen) with $\alpha_{\rm IV} \simeq 2$ ." +CG09 present a general analytical formalism to explain the different power law decays using a fall-back disk. where the variations in à. could potentially be explained by different physics operating within the disk.," CG09 present a general analytical formalism to explain the different power law decays using a fall-back disk, where the variations in $\alpha$ could potentially be explained by different physics operating within the disk." + The results of CGO9 were purely analytical: in this work we present time dependent numerical calculations in order to examine in more detail the potential of the model. and we apply the results to XRT data for one IGRB and one sGRB. taking the best studied of each class.," The results of CG09 were purely analytical; in this work we present time dependent numerical calculations in order to examine in more detail the potential of the model, and we apply the results to XRT data for one lGRB and one sGRB, taking the best studied of each class." + In Section 2 we review the Dainotti relation. an empiricalrelation involving the duration and luminosity of segment IL. in Section 3 we present our detailed numerical," In Section 2 we review the Dainotti relation, an empiricalrelation involving the duration and luminosity of segment II, in Section 3 we present our detailed numerical" +prescut. as well as an intensity miaxinun towards f1ο northeast.,"present, as well as an intensity maximum towards the northeast." + We find tha the axis connecting the centre of the two rines is inclined with respect to the major axis of the houmucuus by either 37 or 58 degrees. dependi1S on which of the two rigsoe is in front. aad assuming au iuclination for t1ο ounluctlus with respect to the pla1ο of the skv of 1) degrees.," We find that the axis connecting the centre of the two rings is inclined with respect to the major axis of the homunculus by either 37 or 58 degrees, depending on which of the two rings is in front, and assuming an inclination for the homunculus with respect to the plane of the sky of 40 degrees." + Lanematics of the gas in the rings is needed to decide which of these two possibilities 1s correct., Kinematics of the gas in the rings is needed to decide which of these two possibilities is correct. + ? show that the iuner honmuculus structure seen at20 µ li represens rogiols ο jucreased counu deusitv., \citetalias{pantin_timmi} show that the inner homunculus structure seen at20 $\mu$ m represents regions of increased column density. + If indeed the riugs trace : cesity cuhanceucut. then they coul be deuser rues 1 la polar nebul:wu similar to the rues in SNI987À (?)..," If indeed the rings trace a density enhancement, then they could be denser rings in a bipolar nebula, similar to the rings in SN1987A \citep{burrows_1987a_hst}." + We also note that this ecomeIV shows strong reseniblaice o that seen iun PNe such as (?) aud (?) including the two riugs. uisalieuimen between the bipolar structure aud the rugs. aud he offset of the ceutral star with respect to rine structure.," We also note that this geometry shows strong resemblance to that seen in PNe such as \citep{sahai_he} and \citep{1999ApJ...522L..69W} + including the two rings, the misalignment between the bipolar structure and the rings, and the offset of the central star with respect to the ring structure." + dto seems reasonable to conclude that oivsieal meclus causing these structures is eeneric a acts in hie[um nass as well as in low mass objects., It seems reasonable to conclude that the physical mechanism causing these structures is generic and acts in high mass as well as in low mass objects. + A nuuDu of mechanisiis to produce double rings iu dipolar nellae have been proposed for PNe (7) axd SNI98TÀ(7).. mt which one applies where has not been established.," A number of mechanisms to produce double rings in bipolar nebulae have been proposed for PNe \citep{icke1988} and \citep{crotts_sn1987a}, but which one applies where has not been established." + T1 all cases the ugs are perpendicular to the uajor axis of 16 nebula. which for jj Car iuplies that the uajor axis of us inner nebula is at a significant angle (37 or 58 degrees) to the major axis of the houmuculus.," In all cases the rings are perpendicular to the major axis of the nebula, which for $\eta$ Car implies that the major axis of this inner nebula is at a significant angle (37 or 58 degrees) to the major axis of the homunculus." + Tt secius cifücul o avoid the couclusiou hat there nust have bee ra change in the orientaion of t1e outflow )etween the nolment of production of f1ο homunculus aud the creaion of the double riuged structure., It seems difficult to avoid the conclusion that there must have been a change in the orientation of the outflow between the moment of production of the homunculus and the creation of the double ringed structure. + This strongly favorrs the binary ποσο]. for he yy Car system., This strongly favours the binary model for the $\eta$ Car system. + The shredded appearance of the skirt 1 itιο DST images aud the proper motion oft16 Condesalons nkicate that the equatorial regions were lughly pertirbed by the ercat eruption., The shredded appearance of the skirt in the HST images and the proper motion of the condensations indicate that the equatorial regions were highly perturbed by the great eruption. + It is therefore likevy that the iies Were produce the exeat. eruption., It is therefore likely that the rings were produced the great eruption. + 1) The chaise of orientation could result fro au asvuinetry in the mass loss duringo the great eruption., 1) The change of orientation could result from an asymmetry in the mass loss during the great eruption. + 2) It could be due to idal interaction of the eccentric |duuarv with material iu its environment., 2) It could be due to tidal interaction of the eccentric binary with material in its environment. + The required lass or such a process can be estimated in the following crude way., The required mass for such a process can be estimated in the following crude way. + A eravitational porurbafion can ac most easilv iu the apoceuter., A gravitational perturbation can act most easily in the apocenter. + Iu a Keplerian motion abou a dunass AL. wit1 eccentricitv e and semi-niajor axis e. the apocenter distance au velocity are given by r£atlde) aud e?2=CMallle)Uboc) .," In a Keplerian motion about a mass $M_{*}$ with eccentricity $e$ and semi-major axis $a$, the apocenter distance and velocity are given by $r=a(1+e)$ and $v^2=\frac{GM_{*}(1-e)}{a(1+e)}$." +Tje required acceleration o change the orbital incination by about | radian is of the order of e/a(l|«j, The required acceleration to change the orbital inclination by about 1 radian is of the order of $v^2/a(1+e)$. + Since the ereat eruptio1 of 1510. about N=160/5.5z29 orbial periods have passed.," Since the great eruption of 1840, about $N=160/5.5 \approx 29$ orbital periods have passed." + T 1rorder to produce the xcquired total change in IN seps. a disturbing mass AL at distauce R would have to fulfil the coudition winAlpay2=ΝΕ1lieΤο}~0051. with AL he reduced mass of the binary.," In order to produce the required total change in $N$ steps, a disturbing mass $\tilde{M}$ at distance R would have to fulfil the condition $\frac{\tilde{M}}{M_{*}}(\frac{a}{R})^2 = +\frac{1}{N}\frac{1-e}{(1+e)^2} \approx 0.0054$, with $M_{*}$ the reduced mass of the binary." + Therefore with €0.6 and Rza. then Afzz5107A...," Therefore with $e=0.6$ and $R\approx a$, then $\tilde{M} \approx 5~10^{-3} \,M_{*}$ ." + A moderate amount of mass close to tle binary could already: be sufficieut to explain the observed change in the svstem orieutation., A moderate amount of mass close to the binary could already be sufficient to explain the observed change in the system orientation. +drives a pressure wave into the jet just upstream of the Mach disk.,drives a pressure wave into the jet just upstream of the Mach disk. + A normal mode of the jet is excited (cL. Harcleeetal. (1998))), A normal mode of the jet is excited (cf. \citet{har}) ) + which steepens to form a conical shock upstream of the Mach disk., which steepens to form a conical shock upstream of the Mach disk. + The influence on the flow is quile significant: constriction of the flow accelerates (he jet upstream of the conical shock to 5~6. but the shock decelerates the flow to +~+: (he corresponding pressure jump is e8.5.," The influence on the flow is quite significant: constriction of the flow accelerates the jet upstream of the conical shock to $\gamma\sim 6$, but the shock decelerates the flow to $\gamma\sim 4$; the corresponding pressure jump is $\sim 8.5$." + In as much as the (tvpes of interaction just discussed are uniquely associated with the leading edge of jets. the structures that result are localized. and will not form a web of propagating shocks along the length of a flow.," In as much as the types of interaction just discussed are uniquely associated with the leading edge of jets, the structures that result are localized, and will not form a web of propagating shocks along the length of a flow." + However. this study highlights that the Mach disk associated with the hotspots of kiloparsec (ancl larger) scale jets mav display orientations revealed by (heir polarization structure Chat bears no simple relation to the jet orientation and accompanied by emission [rom associated oblique shocks. if the head experiences a significantly inhomogeneous ambient medium.," However, this study highlights that the Mach disk associated with the hotspots of kiloparsec (and larger) scale jets may display orientations – revealed by their polarization structure – that bears no simple relation to the jet orientation and accompanied by emission from associated oblique shocks, if the head experiences a significantly inhomogeneous ambient medium." + Further. on anv scale. and in particular on the unresolved sub-parsec scale. à jet that restarts alter a period of quiescence that allows a previously formed channel to relax. or that restarts wilh a new orientation. will display polarized emission (hat is a complex sum of that from (he rotating Mach disk and developing internal shocks. aud which on such scales will evolve within observable time frames.," Further, on any scale, and in particular on the unresolved sub-parsec scale, a jet that restarts after a period of quiescence that allows a previously formed channel to relax, or that restarts with a new orientation, will display polarized emission that is a complex sum of that from the rotating Mach disk and developing internal shocks, and which on such scales will evolve within observable time frames." + 5vxmnmetres seen in (win-jel sources. in particular the S-svimnietric sources. e.g.. (1991).. and the intriguing success of binary black hole models for (he optical-radio waveband behavior of OJ 287. e.g.. ValtonenandLehto(1997).. suggest that jets are subject to a precessional motion. ancl provide a «quantitative explanation lor the possible origin of that motion. respectively.," Symmetries seen in twin-jet sources, in particular the S-symmetric sources, e.g., \citet{lea}, and the intriguing success of binary black hole models for the optical-radio waveband behavior of OJ 287, e.g., \citet{val}, suggest that jets are subject to a precessional motion, and provide a quantitative explanation for the possible origin of that motion, respectively." + Even if precessional motion associated with a massive binary system is nol the cause of a üme-dependent. curvature seen in some parsec-scale flows (e.g.. BL Lac: Alutel et al..," Even if precessional motion associated with a massive binary system is not the cause of a time-dependent curvature seen in some parsec-scale flows (e.g., BL Lac; Mutel et al.," + private communication). an exploration of the impact of precession on flow structure and evolution will give insight into the internal structures to be found in evolving. curved flows.," private communication), an exploration of the impact of precession on flow structure and evolution will give insight into the internal structures to be found in evolving, curved flows." + As a lirst step in exploring the role of precession. we have applied. precession to the inflow of the axisvinmetrie jet discussed above: a jet with inflow Lorentz [actor of 5.," As a first step in exploring the role of precession, we have applied precession to the inflow of the axisymmetric jet discussed above: a jet with inflow Lorentz factor of $5$." + The inflow precesses on a cone of semi-angle 11725 with a lrequency 0.2885 rad measured in {ime units sel bv the inflow radius and speed., The inflow precesses on a cone of semi-angle $11\fdg 25$ with a frequency $0.2885$ rad measured in time units set by the inflow radius and speed. + The relatively large computational domain. ~32x82 jet-radii. meant that we could emplov only three grid levels. with refinement bv x3. providing 11 line cells across the inllowing jet diameter.," The relatively large computational domain, $\sim +32\times 32\times 82$ jet-radii, meant that we could employ only three grid levels, with refinement by $\times 3$, providing $\sim 11$ fine cells across the inflowing jet diameter." + In the approximation that both the inflow speed and bow propagation speed are close to unity. the precession rate implies ~3.75D revolutions of the jet during evolution across the computational domain.," In the approximation that both the inflow speed and bow propagation speed are close to unity, the precession rate implies $\sim 3.75$ revolutions of the jet during evolution across the computational domain." + The, The +The Leading ACDAL cosmology prediets that galaxies form hierarchically in a “bottom. up) fashion (e.g.22). where small perturbations in the dark matter density distribution collapse earlier. than. larger. perturbations. ancl low-mass halos grow by smooth accretion or mergers with other halos. successively building up ever bigger structures.,"The leading $\Lambda$ CDM cosmology predicts that galaxies form hierarchically in a `bottom up' fashion \citep[e.g.][]{WhiteRees1978,GalaxyFormation}, where small perturbations in the dark matter density distribution collapse earlier than larger perturbations, and low-mass halos grow by smooth accretion or mergers with other halos, successively building up ever bigger structures." + But structures falling into bigger svstems during this. process are not always disrupted completely., But structures falling into bigger systems during this process are not always disrupted completely. + As N-body simulations show. the inner cores of infalling objects often. survive the various clisruptive effects. acting on them. Like tidal truncation. tidal shocking or ram-pressure stripping.," As N-body simulations show, the inner cores of infalling objects often survive the various disruptive effects acting on them, like tidal truncation, tidal shocking or ram-pressure stripping." + Lt is believed. that the observed. dwarf galaxies orbiting around the Milkv Way (MW) are examples. of such surviving remnants., It is believed that the observed dwarf galaxies orbiting around the Milky Way (MW) are examples of such surviving remnants. + Based on the first. generation of very high resolution collisionless CDM simulations. 7 and ο pointec oul ἃ very striking apparent cliscrepaney between theoretical predictions for such satellite svstems and— actua observations.," Based on the first generation of very high resolution collisionless CDM simulations, \citet{FirstMSP} and \citet{Moore1999} pointed out a very striking apparent discrepancy between theoretical predictions for such satellite systems and actual observations." + Caüven the very large number of precicte dark matter substructures. there appears to be a dearth of luminous satellites in the Milky Way.," Given the very large number of predicted dark matter substructures, there appears to be a dearth of luminous satellites in the Milky Way." + In fact. the cumulative number of observed satellite 5galaxies. and οἱ Iprecicte substructures above a given. circular velocity. value. cüfferec by à factor of ~10.," In fact, the cumulative number of observed satellite galaxies and of predicted substructures above a given circular velocity value differed by a factor of $\sim 10$." +" Ehis has become known as the ""missing satellite problem'.", This has become known as the `missing satellite problem'. +" "") initial analysis of 2 and ? may have overstated the magnituce of the diserepaney. both because of uncertainties in assigning correct circular velocity values to the observed satellites (2) and because a number of additional faint satellites have been discovered: meanwhile in the MW."," The initial analysis of \citet{Moore1999} and \citet{FirstMSP} may have overstated the magnitude of the discrepancy, both because of uncertainties in assigning correct circular velocity values to the observed satellites \citep{Stoehr2002} and because a number of additional faint satellites have been discovered meanwhile in the MW \citep[see for example][]{SatData9, SatData11, SatData13, SatKinematics, SatData14, +SatData15, SatData16, SatData17, SatData18, SatData19}." + (see that the many low-mass satellites predicted by the N-bocly simulations need to be strongly suppressed. in. luminosity. otherwise a significant cliscrepaney with the observed," However, there is a consensus that the many low-mass satellites predicted by the N-body simulations need to be strongly suppressed in luminosity, otherwise a significant discrepancy with the observed" +assume a constant CR escape time from the Galaxy at higher energies (Ave et al.,assume a constant CR escape time from the Galaxy at higher energies (Ave et al. + 2009)., 2009). +" Under such models. Αν}x££Lr at higher energies while at lower energies N,(£)xEIE7."," Under such models, $N_s(E)\propto E^{-\Gamma}$ at higher energies while at lower energies $N_s(E)\propto E^{-(\Gamma+ a)}$." + The differences we just mentioned are expected in all Kinds of secondary nuclear species like boron. sub-Fe. and anti-protons.," The differences we just mentioned are expected in all kinds of secondary nuclear species like boron, sub-Fe, and anti-protons." + At oresent. data on secondary spectra are available at most only up to ~]00 GeV/n. Future high energy measurements would be crucial © test our model.," At present, data on secondary spectra are available at most only up to $\sim 100$ GeV/n. Future high energy measurements would be crucial to test our model." + In addition. the diffuse ? -ray emission of our Galaxy can also orovide an important check of our model.," In addition, the diffuse $\gamma$ -ray emission of our Galaxy can also provide an important check of our model." +" If the diffuse emission is dominated by the z""-decay 5-ravs. then their intensity would argely follow the proton spectrum at high energies."," If the diffuse emission is dominated by the $\pi^0$ -decay $\gamma$ -rays, then their intensity would largely follow the proton spectrum at high energies." + Therefore. under our model we expect a diffuse spectrum which is steeper han the predictions from other models.," Therefore, under our model we expect a diffuse spectrum which is steeper than the predictions from other models." + In fact. it has already been shown in Yuan et al.," In fact, it has already been shown in Yuan et al." + 2011 that under their model. the > - spectrum would become harder above ~50 GeV. Preliminary results from the FERMI measurements up to ~100 GeV show hat the spectrum is in good agreement with the conventional model assuming a single power-law CR spectrum above a few GeV (Strong 2011).," 2011 that under their model, the $\gamma$ -ray spectrum would become harder above $\sim 50$ GeV. Preliminary results from the FERMI measurements up to $\sim 100$ GeV show that the spectrum is in good agreement with the conventional model assuming a single power-law CR spectrum above a few GeV (Strong 2011)." + The spectrum do show some excess above the model which could well be attributed to unresolved point sources like pulsars., The spectrum do show some excess above the model which could well be attributed to unresolved point sources like pulsars. + Detailed investigation of the diffuse -ray spectrum and also future measurements at even higher energies would be important to check the validity of our model., Detailed investigation of the diffuse $\gamma$ -ray spectrum and also future measurements at even higher energies would be important to check the validity of our model. + In short. we conclude that the apparent change in the spectral index of the CR energy spectra at TeV energies could be a local effect due to nearby SNRs.," In short, we conclude that the apparent change in the spectral index of the CR energy spectra at TeV energies could be a local effect due to nearby SNRs." + A detailed investigation of both the proton and the helium spectra seems to favor this model., A detailed investigation of both the proton and the helium spectra seems to favor this model. + Future measurements of secondary CR spectra and of the Galactic diffuse -ray emission at TeV energies can provide a deeper understanding of the problem., Future measurements of secondary CR spectra and of the Galactic diffuse $\gamma$ -ray emission at TeV energies can provide a deeper understanding of the problem. + Adriani. 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V. $\&$ " +a closer match to the standard Johnson & and V and the Kron- A bandpass.,a closer match to the standard Johnson $B$ and $V$ and the Kron-Cousins $R$ bandpass. + AUX. with its 10221024 TEK CCD. has an unvignetted. circular field diameterof 1.8 arcmin with the pixel size on the sky of 071108 x 071108.," AUX, with its $\times$ 1024 TEK CCD, has an unvignetted, circular field diameter of 1.8 arcmin with the pixel size on the sky of 108 $\times$ 108." + The integration time was 1000 s in the Band V bands each. and 1400 s in the A band.," The integration time was 1000 s in the $B$ and $V$ bands each, and 1400 s in the $R$ band." + The Landolt standard star 94-242 was observed immediately after the primary target at a similar airmass., The Landolt standard star 94-242 was observed immediately after the primary target at a similar airmass. + The weather conditions were very good and stable during the observations., The weather conditions were very good and stable during the observations. + As a result. the FWEM of a stellar object in the images is less than 07775 in all the bands.," As a result, the $FWHM$ of a stellar object in the images is less than 75 in all the bands." + Table | provides a journal of the observations., Table \ref{ObsTab} provides a journal of the observations. + The data were reduced using standard techniques within aandMipAs., The data were reduced using standard techniques within and. +. The frames were debiased and flat tielded using sky flats from the same night., The frames were debiased and flat fielded using sky flats from the same night. + Our astrometry was made using positions of ten reference stars from the USNO-B1.0catalogue!., Our astrometry was made using positions of ten reference stars from the USNO-B1.0. +. We used the ttasks for the astrometric transformation of the images., We used the tasks for the astrometric transformation of the images. + Formal rms errors of the astrometric fit for the RA and DEC were 070066 and 070059 for the R. 070081 and 070035 for the V. and 070079 and 070041 for the B bands respectively. that is better than the pixel size of the images.," Formal rms errors of the astrometric fit for the RA and DEC were 066 and 059 for the R, 081 and 035 for the V, and 079 and 041 for the B bands respectively, that is better than the pixel size of the images." + First of all. the magnitudes of five relatively bright stars visible in the target frame were derived accurately (Figure 1).," First of all, the magnitudes of five relatively bright stars visible in the target frame were derived accurately (Figure \ref{CompStarsFig}) )." + This photometric calibration was carried using the Landolt standard star 94-242., This photometric calibration was carried using the Landolt standard star 94-242. + Our data do not allow us to determine the atmospheric extintion coefficient., Our data do not allow us to determine the atmospheric extintion coefficient. + Instead of this we use the average coethcients ky=0.21]. ky=0.13 and kp=0.09. provided for the Roque de Los Muchachos Observatory. La Palma (Kidgeretal.2003).," Instead of this we use the average coefficients $k_B = 0.21$, $k_V = 0.13$ and $k_R = 0.09$, provided for the Roque de Los Muchachos Observatory, La Palma \citep{Kidger}." +". The signal-to-noise ratios S/N and the magnitude uncertainties Aur were calculated as where /;, is the source flux in counts for a given aperture. e is the gain. 7,, the number of pixels in the source aperture. 7), the number of pixels in area used for the background measurement."," The signal-to-noise ratios $S/N$ and the magnitude uncertainties $\Delta m$ were calculated as where $f_{ap}$ is the source flux in counts for a given aperture, $g$ is the gain, $n_{ap}$ the number of pixels in the source aperture, $n_{bg}$ the number of pixels in area used for the background measurement." +" and c, is the standard deviation of the background in counts (Newberry1991).", and $\sigma_{bg}$ is the standard deviation of the background in counts \citep{Newberry}. +. The stars chosen as secondary standards are marked by numbers in Figure |.. and their magnitudes with errors are listed in Table 2..," The stars chosen as secondary standards are marked by numbers in Figure \ref{CompStarsFig}, and their magnitudes with errors are listed in Table \ref{CompStarsTab}." +" The resulting zero-points for the pulsar frames are B=23797, V=26""06 and A=26220."," The resulting zero-points for the pulsar frames are $B=25 \fm 97$, $V=26 \fm 06$ and $R=26 \fm 20$." + The formal. statistical zero-point errors are less than 00601.," The formal, statistical zero-point errors are less than 01." + However. since only an average curve for atmospheric extinction was used and only one standard star was observed and this only once. the real errors ean be as large as few x 0.01 magnitude.," However, since only an average curve for atmospheric extinction was used and only one standard star was observed and this only once, the real errors can be as large as few $\times$ 0.01 magnitude." + Recent observations of 3C 58 (Shibanovetal.2008). have shown evidence of an optical nebulosity at the same location as the X- counterpart of PSR 02054-6449., Recent observations of 3C 58 \citep{Shibanov08} have shown evidence of an optical nebulosity at the same location as the X-ray counterpart of PSR J0205+6449. + They described it as a faint extended. elliptical structure with no resolved point-like object at its centre.," They described it as a faint extended, elliptical structure with no resolved point-like object at its centre." + Our observations confirm the presence of this nebulosity., Our observations confirm the presence of this nebulosity. + However. its brightest area has a non-elliptical shape withan arc-like filamentary structure.," However, its brightest area has a non-elliptical shape withan arc-like filamentary structure." + Figure 2. (left panel) shows the R band image of this region., Figure \ref{psrFig} (left panel) shows the R band image of this region. + In fact. such a nebula structure is also clearly seen in the V image of 3C 58 presented by Shibanovet in their Figure 3.," In fact, such a nebula structure is also clearly seen in the V image of 3C 58 presented by \citet{Shibanov08} + in their Figure 3." + For consistency with Shibanov et al..," For consistency with Shibanov et al.," + we have performed photometry of this area using the similar elliptical apertures (Table | in their paper), we have performed photometry of this area using the similar elliptical apertures (Table 1 in their paper). + The aperture that formally encapsulates > of the total nebula flux. is shown in Figure 2 (left panel) as a white ellipse.," The aperture that formally encapsulates $\gtrsim$ of the total nebula flux, is shown in Figure \ref{psrFig} (left panel) as a white ellipse." + The measured integral magnitudes of the nebula are B.=2379740.10. V22795+(.05 and Rz22'|52z0.03.," The measured integral magnitudes of the nebula are $B = 23 \fm 97 \pm 0.10$, $V = 22 \fm 95 \pm 0.05$ and $R = 22 \fm 15 \pm 0.03$." + The measured B and V magnitudes agree within errors with those of Shibanovetal.(2008)., The measured B and V magnitudes agree within errors with those of \citet{Shibanov08}. + although our values are a bit brighter., although our values are a bit brighter. + The latter can be explained by non-perfect reduction of the standard stars. as well as by the variable background.," The latter can be explained by non-perfect reduction of the standard stars, as well as by the variable background." + We note however that in the R band the nebula is brighter than the upper limit for the pulsar magnitude taken from Fesenetal.(2008)., We note however that in the R band the nebula is brighter than the upper limit for the pulsar magnitude taken from \citet{Fesen08}. + Furthermore. the R image shows this nebula with some structure significantly above the noise.," Furthermore, the R image shows this nebula with some structure significantly above the noise." + In the right panel of Figure we show the same area of the R image. but with different contrast to emphasise ditterent components of the nebula.," In the right panel of Figure \ref{psrFig} we show the same area of the R image, but with different contrast to emphasise different components of the nebula." + There are three compact objects in the middle of this nebula. marked in Figure 2 as ol. o2 and 03 (the last two could be one are-like object).," There are three compact objects in the middle of this nebula, marked in Figure \ref{psrFig} as o1, o2 and o3 (the last two could be one arc-like object)." + We have measured the magnitudes of these objects. using a 5 pixel radius aperture corresponding to a sky radius of approximately 07554.," We have measured the magnitudes of these objects, using a 5 pixel radius aperture corresponding to a sky radius of approximately 54." + We used the same background as determined previously for the photometry of the nebula., We used the same background as determined previously for the photometry of the nebula. + The measured magnitudes are 22708+0.07 (tol). 2415€0.07 (02) and 2472440.08 (03).," The measured magnitudes are $24 \fm 08 \pm 0.07$ (o1), $24 \fm 15 \pm 0.07$ (o2) and $24 \fm 24 \pm 0.08$ (o3)." + We have also measured the flux from the brightest part of the nebula outside the objects. using the same aperture. and obtained the value z24755x 0.10. giving," We have also measured the flux from the brightest part of the nebula outside the objects, using the same aperture, and obtained the value $\approx 24 \fm 55 \pm 0.10$ , giving" +where! =d/dr.,where $'=d/dr$. +" The eigenmodes have the form B=e“OR,tr)DÀ(cos0) where à—1.2... (=1.2....."," The eigenmodes have the form $B=e^{-t/\tau_{n,l}}F_{n,l}(r)P_l^1(\cos\theta)$ where $n=1,2,...$, $l=1,2,...$." + The fuuction £5; satisfies the equation The boundary conditions are F=Oat r=Oaudr =Rey.," The function $F_ {n,l}$ satisfies the equation The boundary conditions are $F=0$ at $r=0$ and $r=R_{\rm CZ}$." + The lifetimes of various modes. computed for the BP2000 Suu 2001).. are tabulated in Table 1..," The lifetimes of various modes, computed for the BP2000 Sun \citep{BP2000}, are tabulated in Table \ref{table:lifetimes}." +" The n=1. =1 mode has the longest lifetime. 74,4=21Cyr."," The $n=1$, $l=1$ mode has the longest lifetime, $\tau_{1,1}= 24{\rm Gyr}$." + As can be seen from thle table. there are eight modes whose lifetimes exceed the age of the Sun aud three modes whose liletimes exceed 10 billion vears.," As can be seen from the table, there are eight modes whose lifetimes exceed the age of the Sun and three modes whose lifetimes exceed 10 billion years." + Therefore. the toroidal field iu the radiation zone of the Suu cau iu principle have complex structure.," Therefore, the toroidal field in the radiation zone of the Sun can in principle have complex structure." + The profiles Εν) of the eight longest living modes are shown in Figs.," The profiles $F_{n,l}(r)$ of the eight longest living modes are shown in Figs." + 2 aud ὁ and tabulated in Table 2..," \ref{fig:l1} and \ref{fig:l2} + and tabulated in Table \ref{table:Fln}." + AlthoughOm the precediuge calculation does not take into account tliechanges Dm of the solar parameters with time. the ellect of the solar evolution on the lifetimes is expected to be small.," Although the preceding calculation does not take into account the changes of the solar parameters with time, the effect of the solar evolution on the lifetimes is expected to be small." + As a result of hydrogen burning. as the Sun evolves from /=0 to /=LT Gyr. the core contracts aud becomes hotter. while the outer layers of the RZ expaud aud cool.," As a result of hydrogen burning, as the Sun evolves from $t=0$ to $t=4.7$ Gyr, the core contracts and becomes hotter, while the outer layers of the RZ expand and cool." + Since the lifetime depends on the product PT? there is partial cancellation. both in the core aud in the outer part of the RZ. between the e[lects of changiug temperature and clanging size.," Since the lifetime depends on the product $l^2 +T^{3/2}$, there is partial cancellation, both in the core and in the outer part of the RZ, between the effects of changing temperature and changing size." + Moreover. the crossover radius. at which the temperature aud the radius distance o not change. is at about O15. 1991).. close to where the maximum of tlie n=1. /—1 mode.," Moreover, the crossover radius, at which the temperature and the radius distance do not change, is at about $0.15 +R_\odot$ \citep{Demarque91}, close to where the maximum of the $n=1$, $l=1$ mode." + Iu this Section we discuss observational upper bouuds on the allowed streugth of the magnetic field in the radiation zone., In this Section we discuss observational upper bounds on the allowed strength of the magnetic field in the radiation zone. + A variety of independent. arguments rule out fields in excess of <10? G. We briefly παίσο the arguments [rom helioseismological measurements of tlie souud speed and from the measurements of the *B neutrino flux., A variety of independent arguments rule out fields in excess of $\lesssim 10^8$ G. We briefly summarize the arguments from helioseismological measurements of the sound speed and from the measurements of the $^8$ B neutrino flux. + A much stronger bound on the axisyiunetrical fields comes from the measured oblateuess of the Suu., A much stronger bound on the axisymmetrical fields comes from the measured oblateness of the Sun. + We derive a bound ou the strenetl of the field specializing to the longest living 1»=1. /=1 mocle.," We derive a bound on the strength of the field specializing to the longest living $n=1$, $l=1$ mode." + We also point out that a comparable bound could be obtained from the helioseisinological measurements of the mode splitting., We also point out that a comparable bound could be obtained from the helioseismological measurements of the mode splitting. +" The magnetic field of strength B adds a contribution to pressure pj,=B?/8s. correspondingly decreasing the gas pressure py."," The magnetic field of strength $B$ adds a contribution to pressure $p_m=B^2/8\pi$, correspondingly decreasing the gas pressure $p_{\rm g}$." + By estimating the change in the gas temperature as [ὁ11[ο/papy. oue obtaius a bound ou the allowed strength of the field in the core aud iu the RZ from the observed [lux of the B neutrinos aud the helioseismological measurements of the sound speed.," By estimating the change in the gas temperature as $|\delta T/T| \sim p_m/p_{\rm g}$, one obtains a bound on the allowed strength of the field in the core and in the RZ from the observed flux of the $^8$ B neutrinos and the helioseismological measurements of the sound speed." + The two constraints⋅ are complimentary.⋅ because tlie x B neutrinos are produced very close to the ceuter of," The two constraints are complimentary, because the $^8$ B neutrinos are produced very close to the center of" +The simulations have been performed with the Tree-SPLI code VINE.,The simulations have been performed with the Tree-SPH code VINE. + A full description of VINE can be found in ? and ?.., A full description of VINE can be found in \citet{Wetzstein2008} and \citet{Nelson2008}. + VINE is parallelised with OpenAL? directives., VINE is parallelised with OpenMP directives. + It invokes a leapfrog integrator. and individual particle time steps with a CEL tolerance parameter of 0.1.," It invokes a leapfrog integrator, and individual particle time steps with a CFL tolerance parameter of 0.1." + Gravitational accelerations are estimated using a tree. with opening angle 6=0.005 (?).," Gravitational accelerations are estimated using a tree, with opening angle $\theta=0.005$ \citep{Springel2001}." +" The gravitational softening length equals the hydrodvnamical smoothing leneth (2)... which is adapted so that cach particle has S,=50x20 neighbours."," The gravitational softening length equals the hydrodynamical smoothing length \citep{BateBurkert97}, which is adapted so that each particle has ${\cal N}_{_{\rm NEIB}}=50\pm20$ neighbours." + Hydrodynamical forces are treated with periodic boundary conditions. but gravitational forces are not.," Hydrodynamical forces are treated with periodic boundary conditions, but gravitational forces are not." +" Artificial viscosityis treated using the standard prescription of ? with a,= band Z4.=2. plus the Dalsara switch (?).."," Artificial viscosityis treated using the standard prescription of \citet{Gingold1983} with $\alpha_{_{\rm AV}}=1$ and $\beta_{_{\rm AV}}=2$, plus the Balsara switch \citep{Balsara1995}." + We have experimented with time-varving viscosity (?).. but find that it makes no significant dillerence.," We have experimented with time-varying viscosity \citep{Morris1997}, but find that it makes no significant difference." +" For the purpose of calculating the gross thermodynamics. we assume that the gas is pure molecular hydrogen. with ratio of specific heats ,=7/5 and molecular weight f=2: thus for simplicity we are neglecting the Lact that the rotational degrees of freedom of LL, are not Lully excited at. low temperatures. and. we are neglecting the contribution of helium top."," For the purpose of calculating the gross thermodynamics, we assume that the gas is pure molecular hydrogen, with ratio of specific heats $\gamma=7/5$ and molecular weight $\mu=2$; thus for simplicity we are neglecting the fact that the rotational degrees of freedom of $_{_2}$ are not fully excited at low temperatures, and we are neglecting the contribution of helium to $\mu$." + The equation of state is and the isothermal sound speed is In solving the energy equation we adopt the following procedure., The equation of state is and the isothermal sound speed is In solving the energy equation we adopt the following procedure. +" At low densities.""A p«pou,=1013""gemH"". we use the cooling rates computed by 2.. with the constraint that the temperature is not allowed to fall below 2=91x."," At low densities, $\rho<\rho_{_{\rm CRIT}}=10^{-13}\,{\rm g}\,{\rm cm}^{-3}$, we use the cooling rates computed by \citet{Neufeld1995}, with the constraint that the temperature is not allowed to fall below $T=9\,{\rm K}$." +" At high densities. pp,quz. we switch olf the radiative cooling.and the gas then evolves adiabaticallv."," At high densities, $\rho>\rho_{_{\rm CRIT}}$, we switch off the radiative cooling,and the gas then evolves adiabatically." +" ? have computed the equilibrium abundances of kev molecules and atoms (CO. 11,O. LL,. LICL O,. C and O). and the resulting net cooling rate due to line emission. A. as a function of density. and temperature. using the local velocity. eradicnt method to treat optical-depth ellects."," \citet{Neufeld1995} have computed the equilibrium abundances of key molecules and atoms (CO, $_{_2}$ O, $_{_2}$, HCl, $_{_2}$, C and O), and the resulting net cooling rate due to line emission, $\Lambda$, as a function of density and temperature, using the local velocity gradient method to treat optical-depth effects." +" Neufeld has kindly supplied these rates in the form of a look-up table for densities in. the range 107(ns)cnm7)107"" and temperatures. in the range 10x(ZI)2500.", Neufeld has kindly supplied these rates in the form of a look-up table for densities in the range $10^3\le\left(n(H_2)/{\rm cm}^{-3}\right)\le 10^{10}$ and temperatures in the range $10\le\left(T/{\rm K}\right)\le 2500$. + Most. previous work in this Ποιά has used a barotropic equation of state and. therefore does not treat thermal inertia elfects (e.g.2?7).. although 7 use the same procedure as us.," Most previous work in this field has used a barotropic equation of state and therefore does not treat thermal inertia effects \citep[e.g.][]{Goodwin2004, Bate1998, Matsumoto2003}, although \citet{Banerjee2004} use the same procedure as us." + By invoking line cooling we hope to capture more realistically the thermal behaviour of low-clensity material as it falls onto the protostellar disc and then spirals into the central protostar., By invoking line cooling we hope to capture more realistically the thermal behaviour of low-density material as it falls onto the protostellar disc and then spirals into the central protostar. + Ideally: racliative cooling should also take account of continuum emission from dust. in the regime where the gas and. dust. are thermally. coupled. for exemple by using the algorithm developed by ?..," Ideally radiative cooling should also take account of continuum emission from dust, in the regime where the gas and dust are thermally coupled, for example by using the algorithm developed by \citet{Stamatellos2007}." + Unfortunately the computational machinery to treat both the line cooling regime and the dust cooling regime does not vet exist., Unfortunately the computational machinery to treat both the line cooling regime and the dust cooling regime does not yet exist. + We are currently. developing an algorithm to combine these regimes., We are currently developing an algorithm to combine these regimes. +" The core is modeled with 430.000 SPIE particles. cach raving mass ny,=Lt.10*M.."," The core is modeled with 430,000 SPH particles, each having mass $m_{_{\rm SPH}}=1.4\times 10^{-5}\,{\rm M}_\odot$." +" Following ?.. the minimum resolvable mass is therefore 2N4,4,=1410 M..."," Following \citet{BateBurkert97}, the minimum resolvable mass is therefore $2{\cal N}_{_{\rm NEIB}}m_{_{\rm SPH}}=1.4\times 10^{-3}\,{\rm M}_\odot$ ." +" For comparison the mininium possible Jeans mass (corresponding to p—cay, and P=9I) is L5.10PAL."," For comparison the minimum possible Jeans mass (corresponding to $\rho=\rho_{_{\rm CRIT}}$ and $T=9\,{\rm K}$ ) is $1.5\times 10^{-3}\,{\rm M}_\odot$." + ''herefore the Jeans mass is always resolved., Therefore the Jeans mass is always resolved. + In. adcdition here are 24.000 ambient particles exerting the external »essure that contains the core (see Eqn. 3)).," In addition there are 24,000 ambient particles exerting the external pressure that contains the core (see Eqn. \ref{EQN:PEXT}) )." +" These particles also have mass m4,=ld10M.. but they have much higher temperature. and their density is adjusted. so hat they fill the space between the core and the periodic rotunclaries."," These particles also have mass $m_{_{\rm SPH}}=1.4\times 10^{-5}\,{\rm M}_\odot$, but they have much higher temperature, and their density is adjusted so that they fill the space between the core and the periodic boundaries." + We do not use sink particles., We do not use sink particles. + Instead. we Limit he CPU time by imposing a minimum smoothing length. hus=2AU.," Instead we limit the CPU time by imposing a minimum smoothing length, $h_{_{\rm MIN}}=2\,{\rm AU}$." + Consequently we are able to resolve the details of the protostellar disc. both racially and vertically. but not he central protostar.," Consequently we are able to resolve the details of the protostellar disc, both radially and vertically, but not the central protostar." +" Any SPILL particle whose density. rises above ϱ=101Lgem"" is: presumed to be part of⋅ a protostar."," Any SPH particle whose density rises above $\rho_\star\!=\!10^{-11}\,{\rm g}\,{\rm cm}^{-3}$ is presumed to be part of a protostar." +" MThe primary protostar is the one that forms first. at the centre of the core. and /,,o is its time of formation."," The primary protostar is the one that forms first, at the centre of the core, and $t_{_{\rm O}}$ is its time of formation." +" Values of ἐς, are given in Table 1...", Values of $t_{_{\rm O}}$ are given in Table \ref{TAB:SIMUS}. +" ἐς, increases monotonically with (3,. because more rapidly rotating cores take longer to reach the critical density. as."," $\;t_{_{\rm O}}$ increases monotonically with $\Omega_{_{\rm O}}$, because more rapidly rotating cores take longer to reach the critical density, $\rho_\star$." +" We define the disc to be all material with density we]in the range 10.1σοι3xp10""mecm5°. since in general this material has [o]29|o] Qvhere ος, and ον are he azimuthal and racial components of velocity)"," We define the disc to be all material with density in the range $10^{-16}\,{\rm g}\,{\rm cm}^{-3}\leq\rho\leq 10^{-11}\,{\rm g}\,{\rm cm}^{-3}$, since in general this material has $|v_\phi| \gg |v_r|$ (where $v_\phi$ and $v_r$ are the azimuthal and radial components of velocity)." + The evolution of a core can usefully be divided. into wo distinct phases., The evolution of a core can usefully be divided into two distinct phases. + During the radiative cooling is very cllicient anc theeas temperature is roughly constant., During the radiative cooling is very efficient and thegas temperature is roughly constant. + Vhroughout most of the isothermal hase centrifugal acceleration is a secondary elfect. and he core collapses on an approximately [reefall time scale.," Throughout most of the isothermal phase centrifugal acceleration is a secondary effect, and the core collapses on an approximately freefall time scale." + The starts as soon as a disc forms., The starts as soon as a disc forms. + The xttern of evolution during the Protostellar Phase depends critically on QO Beye , The pattern of evolution during the Protostellar Phase depends critically on $\Omega_{_{\rm O}}$ . +"In cores with high ο. (a) formation of the »imary protostar is delaved until a substantial,,. dise has built"," In cores with high $\Omega_{_{\rm O}}$ , (a) formation of the primary protostar is delayed until a substantial disc has built" +of Bildstenοἱal.(1997). ideal gas equation of state. ancl zero metallicity opacities) have 1-0 deviations [rom our reference caleulation that are of order124%.,"of \citet{bil97}, ideal gas equation of state, and zero metallicity opacities) have $\sigma$ deviations from our reference calculation that are of order." +.. Thus. unless we are ignoring input physics that impact stellar structure of order the error associated with treating the eas physics as an ideal gas. our errors are accurate and are not underestimates.," Thus, unless we are ignoring input physics that impact stellar structure of order the error associated with treating the gas physics as an ideal gas, our errors are accurate and are not underestimates." + Transformation from a theoretical luminosity or effective temperature (o an observable magnitude svstem requires accurate knowledge of the bolometric correction., Transformation from a theoretical luminosity or effective temperature to an observable magnitude system requires accurate knowledge of the bolometric correction. + A quantitative discussion on the impact of the bolometric correction uncertaintv on the derived LDB ages does not exist in (he literature., A quantitative discussion on the impact of the bolometric correction uncertainty on the derived LDB ages does not exist in the literature. + Thus. in (his section we briefly characterize this additional error source.," Thus, in this section we briefly characterize this additional error source." + There are three methods for bolometric correction determination: empirical. (heoretical. and mixed (theoretical method with empirical constraints).," There are three methods for bolometric correction determination: empirical, theoretical, and mixed (theoretical method with empirical constraints)." + Observations for stus of known distance and wide spectral coverage allow an empirical relation between the bolometric magnitude of a star and its broad-band colors (Bessell1991:Monetetal.1992).," Observations for stars of known distance and wide spectral coverage allow an empirical relation between the bolometric magnitude of a star and its broad-band colors \citep{bes91,mon92}." + With improved atmospheres. fully (theoretical caleulations for the bolometric corrections are possible (ILauschildt.Allard.&Baron 1999).," With improved atmospheres, fully theoretical calculations for the bolometric corrections are possible \citep{hau99}." +. A mixture of the theoretical ancl empirical methods attempts to improve the coverage of bolometric corrections where observations are sparse. by collecting theoretical bolometric calculations and then correcting (hem enipirically (Lejeune.Cuisinier.&Buser1997).," A mixture of the theoretical and empirical methods attempts to improve the coverage of bolometric corrections where observations are sparse, by collecting theoretical bolometric calculations and then correcting them empirically \citep{lej97}." +. For our reference bolometric corrections. we choose the theoretical bolometric corrections ol Hatsehildt.Allard.&Baron(1999).," For our reference bolometric corrections, we choose the theoretical bolometric corrections of \citet{hau99}." +. At fixed elfective temperature. the top panel in Figure 9 shows the comparison between our reference bolometric corrections aud [ος other bolometrie correction determinations.," At fixed effective temperature, the top panel in Figure \ref{bcerr} shows the comparison between our reference bolometric corrections and four other bolometric correction determinations." + The solid line represents (he comparison with the mixed theorv/empirical calculations of Lejeune.Cuisinier.&Buser (1997).., The solid line represents the comparison with the mixed theory/empirical calculations of \citet{lej97}. . + The difference is in (he sense of Lejeune.Cuisinier.&Buser(1997) minus the bolometric corrections ol Iauschildt.Allard.&Baron(1999)., The difference is in the sense of \citet{lej97} minus the bolometric corrections of \citet{hau99}. +.. The long-dash line shows the comparison with the theoretical models used by Lejeune.Cuisinier.&Buser(1997) before applving the empirical corrections., The long-dash line shows the comparison with the theoretical models used by \citet{lej97} before applying the empirical corrections. + The dash-dot line ancl short-dash lines compare the results to empirical relations of Monetetal.(1992) and Bessell(1991).. respectively.," The dash-dot line and short-dash lines compare the results to empirical relations of \citet{mon92} and \citet{bes91}, respectively." + The empirical corrections. in practice. are given as a function of observed photometric color. aud blackhocly spectral enerev distributions fit to the stellar spectra provide the effective temperature scale.," The empirical corrections, in practice, are given as a function of observed photometric color, and blackbody spectral energy distributions fit to the stellar spectra provide the effective temperature scale." + Due to the departure from a blackbody. the blackbody fitting method for determiningthe effective temperature becomes increasingly inaccurate and highlv questionable lor temperatures.< 3000 IX (Monetetal. 1992)..," Due to the departure from a blackbody, the blackbody fitting method for determiningthe effective temperature becomes increasingly inaccurate and highly questionable for temperatures$<$ 3000 K \citep{mon92}. ." +"Laboratory emission spectra of hot NIL, were recorded using a svstem simular to that used by Nassar&Bernath(2003) to study. hot Cll).",Laboratory emission spectra of hot $_{3}$ were recorded using a system similar to that used by \citet{nassar03} to study hot $_{4}$. + A diagram of the experiment is shown in Figure 2.., A diagram of the experiment is shown in Figure \ref{fig2}. +" An alumina (AI50,) sample tube is sealed with windows ancl evacuated.", An alumina $_{2}$ $_{3}$ ) sample tube is sealed with windows and evacuated. + A constant slow flow of δω gas is introduced to the svstem and a constant pressure is maintained using a needle valve which helps to minimize (he build up of impuriües aud loss of sample within the svstem., A constant slow flow of $_{3}$ gas is introduced to the system and a constant pressure is maintained using a needle valve which helps to minimize the build up of impurities and loss of sample within the system. + Surrounding the central 51 em of the AlSO4 tube (121 cm long) is a controllable tube furnace capable of maintaining stable temperatures up to 13707 C. The radiation exiting the AlsO; tube is focussed into a Fourier transform infrared (FT-IR) spectrometer using a lens aud the clistance between (he window on (he end of the Al3O4 tube and the FT-IR. spectrometer entrance aperture was 20 cm., Surrounding the central 51 cm of the $_{2}$ $_{3}$ tube (121 cm long) is a controllable tube furnace capable of maintaining stable temperatures up to $^{\circ}$ C. The radiation exiting the $_{2}$ $_{3}$ tube is focussed into a Fourier transform infrared (FT-IR) spectrometer using a lens and the distance between the window on the end of the $_{2}$ $_{3}$ tube and the FT-IR spectrometer entrance aperture was 20 cm. + This volume was purged with ‘drv air (11Ο absent) to minimize the presence of H3O absorption lines in the enission spectra., This volume was purged with `dry air' $_{2}$ O absent) to minimize the presence of $_{2}$ O absorption lines in the emission spectra. +Discovered in 2006 and first observed to be active on UT 2011 Aug. 30 (Ilsieh οἱ al.,Discovered in 2006 and first observed to be active on UT 2011 Aug. 30 (Hsieh et al. + 201141. little is vet known about this object.," 2011d), little is yet known about this object." + With reported absolute magnitude // = 16.6 and an assumed albedo py = 0.04. appropriate to its outer belt location. (he estimated diameter is ~3 km.," With reported absolute magnitude $H$ = 16.6 and an assumed albedo $p_V$ = 0.04, appropriate to its outer belt location, the estimated diameter is $\sim$ 3 km." + Two thin tails. one near the projected orbit ancl another roughly antisolar. show that dust leaves (he nucleus very slowly and point to mass loss over a protracted period.," Two thin tails, one near the projected orbit and another roughly antisolar, show that dust leaves the nucleus very slowly and point to mass loss over a protracted period." + There is no reported evidence for gas or for repetitive mass loss. but new observations of this object are still be acquired at the time of writing.," There is no reported evidence for gas or for repetitive mass loss, but new observations of this object are still be acquired at the time of writing." +" The appearance is (vpically that of a point-like nucleus with a thin tail (or ""trail) οἱ 10 yan sized dust following in the projected orbit.", The appearance is typically that of a point-like nucleus with a thin tail (or “trail”) of 10 $\mu$ m sized dust following in the projected orbit. + Order of magnitude mass loss rates in eas. inferred. [rom spectroscopy. (Licandro et al.," Order of magnitude mass loss rates in gas, inferred from spectroscopy (Licandro et al." + 2011) and m dust. inferred. [rom surface photometry of the tail (Hsieh et al.," 2011) and in dust, inferred from surface photometry of the tail (Hsieh et al." + 2004) reach 70.02 kg ., 2004) reach $\sim$ 0.02 kg $^{-1}$. + The thin tail indicates (hat particles are ejected from the nucleus al speeds ~1.5 5 ms 1. comparable to the escape speed from the 3.80.6 km diameter nucleus., The thin tail indicates that particles are ejected from the nucleus at speeds $\sim$ 1.5 m $^{-1}$ comparable to the escape speed from the $\pm$ 0.6 km diameter nucleus. + A cometary appearance was detected [or 176P/LINEAR. over a single. month-long interval in 2005 n(Ilseh et al.," A cometary appearance was detected for 176P/LINEAR over a single, month-long interval in 2005 (Hsieh et al." + --2011a)., 2011a). + During this time. the object showed a fan-shaped tail and was than the bare nucleus. leading to an implied dust mass 10? ke.," During this time, the object showed a fan-shaped tail and was about brighter than the bare nucleus, leading to an implied dust mass $\sim$ $^5$ kg." + The properties of the dust can be approximately matched by models in which the characteristic particle size is 10 jam. the ejection speed ~5 ms + and the dust production rale ~ 0.07 kg I. all similar to values inferred in 1990. The 4.00.4 km diameter nucleus rotales with a period near 22.2 hr (Table 2)).," The properties of the dust can be approximately matched by models in which the characteristic particle size is 10 $\mu$ m, the ejection speed $\sim$ 5 m $^{-1}$ and the dust production rate $\sim$ 0.07 kg $^{-1}$, all similar to values inferred in 133P. The $\pm$ 0.4 km diameter nucleus rotates with a period near 22.2 hr (Table \ref{physical}) )." + Like P/2008 R1. the nucleus of 238P/Read is ny. with a diameter 70.5 km (118191 et al.," Like P/2008 R1, the nucleus of 238P/Read is tiny, with a diameter $\sim$ 0.8 km (Hsieh et al." + 2011b)., 2011b). + It was observed (ο be in an active state in both 2005 and 2010. but not in between.with a coma dust mass of order 10? ke and a production rate estimated (Irom published photometry) near ~0.1 ke !.," It was observed to be in an active state in both 2005 and 2010, but not in between,with a coma dust mass of order $^5$ kg and a production rate estimated (from published photometry) near $\sim$ 0.1 kg $^{-1}$." + Also like P/2008 RI. 238P/Reacl is dynamically short-lived. with a survival time of order 20 Myr (HIaghighipour 2009).," Also like P/2008 R1, 238P/Read is dynamically short-lived, with a survival time of order 20 Myr (Haghighipour 2009)." + The object was observed to be active [rom2010 September (to2011 January. at 2 — 2.6 io 2.7 AU.," The object was observed to be active from 2010 September to 2011 January, at $R$ = 2.6 to 2.7 AU." + Moreno et al. (, Moreno et al. ( +2011b) inferred. dust production at {he peak rate of ~4 kg ! with centimeter-sized particles ejected at about0.1 to0.2 ms .,"2011b) inferred dust production at the peak rate of $\sim$ 4 kg $^{-1}$ ,with centimeter-sized particles ejected at about 0.1 to 0.2 m $^{-1}$ ." + A limit to the outgassing, A limit to the outgassing +DTD is that it follows a power law: it decreases inversely proportional to time (galaxy age).,DTD is that it follows a power law: it decreases inversely proportional to time (galaxy age). +" The only caveat is that the observations are presented per unit delay time per century for a single starburst population whose total K-band luminosity Zao at an age of I] Gyr is 10"" Εκ... while all population synthesis models give SN Ia rates 1n number of events per unit time for a single starburst with a given total initial mass."," The only caveat is that the observations are presented per unit delay time per century for a single starburst population whose total K-band luminosity $L_{K,0}$ at an age of 11 Gyr is $10^{10}$ $_{\mathrm{K},\odot}$, while all population synthesis models give SN Ia rates in number of events per unit time for a single starburst with a given total initial mass." + This implies that a conversio1 is required between SN Ia rate per K-band luminosity (ΝΙΚ) and SN Ia rate per mass (SNuM)., This implies that a conversion is required between SN Ia rate per K-band luminosity (SNuK) and SN Ia rate per mass (SNuM). + Totanietal.(2008) adopt a conversion factor obtained from a spectral energy distributio template for an exponentially decaying star formation rate with a timescale of 0.1 Gyr., \citet{totani2008} adopt a conversion factor obtained from a spectral energy distribution template for an exponentially decaying star formation rate with a timescale of 0.1 Gyr. + This factor depends on the adopted metallicity Z and initial mass function (IMF). anc values are given for different choices.," This factor depends on the adopted metallicity $Z$ and initial mass function (IMF), and values are given for different choices." + The cases assuming solar metallicity and the Chabrier(2003) IMF best match our use (see Sect., The cases assuming solar metallicity and the \citet{chabrier2003} IMF best match our use (see Sect. + 2) of Z=0.02 and the Kroupaetal.(1993) IMF., 2) of $Z=0.02$ and the \citet{kroupa1993} IMF. +" Following personal communication with Totant. we take à value based on a combination of these models: In this equation. M, is defined as the integrated star formation rate up to the age of the galaxy. which in the case of a single starburst is nothing else than the total mitial mass."," Following personal communication with Totani, we take a value based on a combination of these models: In this equation, $M_*$ is defined as the integrated star formation rate up to the age of the galaxy, which in the case of a single starburst is nothing else than the total initial mass." + Since this conversion factor is taken to be constant for the entire redshift range. any change (or uncertainty) in it will only influence the absolute values of the SN Ia rate in SNuM. and not the morphological shape of the DTD in those units.," Since this conversion factor is taken to be constant for the entire redshift range, any change (or uncertainty) in it will only influence the absolute values of the SN Ia rate in SNuM, and not the morphological shape of the DTD in those units." + The question addressed in this work will thus pertain to the dominant formation scenario in the case of elliptical galaxies., The question addressed in this work will thus pertain to the dominant formation scenario in the case of elliptical galaxies. + As a novelty. we will specifically investigate the influence of the description of mass transfer efficiency during Roche lobe overflow (RLOF).," As a novelty, we will specifically investigate the influence of the description of mass transfer efficiency during Roche lobe overflow (RLOF)." + In Sect., In Sect. + 2. we give a brief description of our population number synthesis code. and we elaborate on the new physical effects which have been incorporated as well as on the details of systems leading to a SN Ia. Sect.," 2, we give a brief description of our population number synthesis code, and we elaborate on the new physical effects which have been incorporated as well as on the details of systems leading to a SN Ia. Sect." + 3 contains our new results on the progenitors of SNe la. as well as a short elaboration on a special kind of merger.," 3 contains our new results on the progenitors of SNe Ia, as well as a short elaboration on a special kind of merger." + In Sect., In Sect. + 4. a comparison is made between the results of our code and those published by groups undertaking similar studies.," 4, a comparison is made between the results of our code and those published by groups undertaking similar studies." + For our study. we use the Brussels population number synthesis code developed over the years at the Astrophysical Institute of the Vrije Universiteit Brussel.," For our study, we use the Brussels population number synthesis code developed over the years at the Astrophysical Institute of the Vrije Universiteit Brussel." + A very extended review of the elements contained in this code is given by DeDonder&Van-beveren (2004)., A very extended review of the elements contained in this code is given by \citet{dedonder2004}. +. We repeat here only the most important points and the changes which have been implemented since then. reflecting recent developments in the descriptio of certain physical processes.," We repeat here only the most important points and the changes which have been implemented since then, reflecting recent developments in the description of certain physical processes." + Our code starts from an instantaneous starburst normalized to 10° M. with Z=0.02 and a binary fraction of100%.., Our code starts from an instantaneous starburst normalized to $10^6$ $_{\odot}$ with $Z=0.02$ and a binary fraction of. + By default the stars are drawn from a Kroupaetal.(1993) IMF. flat mass ratio distribution and Abt(1983) initial orbital separation distribution.," By default the stars are drawn from a \citet{kroupa1993} IMF, flat mass ratio distribution and \citet{abt1983} initial orbital separation distribution." + The implications of other choices will be discussed in Sect., The implications of other choices will be discussed in Sect. + 3., 3. + Supernova rates are given in the output as number of events per Gyr as a function of time elapsed since starburst., Supernova rates are given in the output as number of events per Gyr as a function of time elapsed since starburst. + As stated above. a main new factor in this investigation is the influence of the description of the RLOF.," As stated above, a main new factor in this investigation is the influence of the description of the RLOF." + Whether or not this process Is conservative Is indicated by the parameter f., Whether or not this process is conservative is indicated by the parameter $\beta$. + In the conservative case. 6=|. the total mass and angular momentum of the system are conserved during RLOF.," In the conservative case, $\beta = 1$, the total mass and angular momentum of the system are conserved during RLOF." + Matter is transferred from the initially more massive star (Mj) to the initially less massive star (M). whereby the latter accretes all the matter lost by the donor.," Matter is transferred from the initially more massive star $M_1$ ) to the initially less massive star $M_2$ ), whereby the latter accretes all the matter lost by the donor." + This process can continue after the donor star has become less massive than the accretor star. a phase known as mass ratio reversal.," This process can continue after the donor star has become less massive than the accretor star, a phase known as mass ratio reversal." + Angular momentum ts exchanged between the two components. and with the orbit.," Angular momentum is exchanged between the two components, and with the orbit." + The resulting orbital period is given by The indices in this formula. 1 for initial and f for final. are with respect to the RLOF phase.," The resulting orbital period is given by The indices in this formula, i for initial and f for final, are with respect to the RLOF phase." + This ratio is smaller or larger than unity depending on whether mass transfer continued until after mass ratio reversal., This ratio is smaller or larger than unity depending on whether mass transfer continued until after mass ratio reversal. + When mass transfer is non-conservative. the following relation holds: P is thus/4/ the fraction of material lost by the donor which ts accepted by the aecretor.," When mass transfer is non-conservative, the following relation holds: $\beta$ is thus the fraction of material lost by the donor which is accepted by the accretor." + The amount of mass lost from the system is readily found by integration over the time during which mass loss takes place., The amount of mass lost from the system is readily found by integration over the time during which mass loss takes place. + To determine the amount of angular momentum lost. and thus the final orbital period. however. requires knowledge of the way in which the non-acereted material leaves the system.," To determine the amount of angular momentum lost, and thus the final orbital period, however, requires knowledge of the way in which the non-accreted material leaves the system." + The two most common approaches. which will be considered in this work. ts that either the matter is lost with the specific angular momentum of the accretor. or with the orbital angular momentum of a corotating point at the second Lagrangian point (L>).," The two most common approaches, which will be considered in this work, is that either the matter is lost with the specific angular momentum of the accretor, or with the orbital angular momentum of a corotating point at the second Lagrangian point $\mathrm{L}_2$ )." + The first assumption is justified rf the mass loss can be treated as an enhanced stellar wind. or at least as a process which is symmetrical in the equatorial plane of the accretor.," The first assumption is justified if the mass loss can be treated as an enhanced stellar wind, or at least as a process which is symmetrical in the equatorial plane of the accretor." + There is however no physical mechanism which is known to make this possible in normal intermediate mass binaries., There is however no physical mechanism which is known to make this possible in normal intermediate mass binaries. + The formulas thus obtained for orbital period variation during non-conservative RLOF are the following Donder&Vanbeveren 2004)., The formulas thus obtained for orbital period variation during non-conservative RLOF are the following \citep[see e.g.][]{dedonder2004}. +. ForO01 aud (2) the Iartree xteutial is orbital iudepenudeut., We note that (1) $\beta_Z > 1$ and (2) the Hartree potential is orbital independent. + The former implies the occupation of ο20 Landau levels is small compared o the »=0 state., The former implies the occupation of $n > 0$ Landau levels is small compared to the $n=0$ state. +" Therefore equation (6)) is amenable o a solution in which we break the sinele-orbital Tamiltonian (5/70) into a Oth order term (003 and a Ist order term (Py, ", Therefore equation \ref{5.4}) ) is amenable to a solution in which we break the single-orbital Hamiltonian $h^{(0)}$ ) into a 0th order term $\tilde h^{(0)}$ ) and a 1st order term $\tilde h^{(1)}$ ). +The 100 term will represent a o»urelv sinegle-configurational calculation for the sinele-orbital wavefunction., The $\tilde h^{(0)}$ term will represent a purely single-configurational calculation for the single-orbital wavefunction. +" The orbitalindependence of the Taitree-potential will eusurethat we can restrict the iterative calculation of the Hartree wavefunction to the n— Osinele-configurational (one dimeusional) calculation: ounce we have determined the longitudinal part of this wavefunction (fir) iteratively, it will be inuuecliately applicable to the construction of the multi-confieurational ILutree sinele-orbital wavefunctions without further uecd for iterative calculations."," The orbital-independence of the Hartree-potential will ensurethat we can restrict the iterative calculation of the Hartree wavefunction to the $n=0$ single-configurational (one dimensional) calculation; once we have determined the longitudinal part of this wavefunction $f_{m\nu}$ ) iteratively, it will be immediately applicable to the construction of the multi-configurational Hartree single-orbital wavefunctions without further need for iterative calculations." +" This represents zn enormous conrputational simplification over the schemes emploved for hwdrosen atoms (Rodsucrctal.19814:Forsterot1981) iu which the longitudinal wavefuuctiou for cach Laudau level in the multi-configurational solution iust be calculated separately,"," This represents an enormous computational simplification over the schemes employed for hydrogen atoms \citep{roesner84, forster84} in which the longitudinal wavefunction for each Landau level in the multi-configurational solution must be calculated separately." +" With a set of cigenfiuctious (X,,55) for the Oth order siugle-coufigurational ILuuiltouiau 5/9! determined we can solve for the imulti-coufieurational siugle-orbital ]: . -- ↖↖↽⋜↧↖↽↸∖↕∏⊔⊳⊓∪∐↴∖↴≺∩⊔⊔∕⋝⋜⋯≺↧↸∖∐↸∖↥⋅∶↴∙⋅↖↽↸∖↕∶↴∙↸∖∐↖↽⋜↧↕⋯∖↴∖↴≺←⊔⊔∕⋝∏↴∖↴∐↓∶↴∙AU) ∙↜ perturbationfaati theory ow]with jl)D.", With a set of eigenfunctions $\chi_{nm\nu}$ ) for the 0th order single-configurational Hamiltonian $\tilde h^{(0)}$ determined we can solve for the multi-configurational single-orbital wavefunctions $\phi_{m\nu}$ ) and energy eigenvalues $\epsilon^{(0)}_{m\nu}$ ) using perturbation theory with $\tilde h^{(1)}$. +" We7 constructconstinc the total; wavefunction. Wir) for II"" bw forming a completely antisviunetized spatial wavefunction from the ον (as required by the FSP assumption).", We construct the total wavefunction $\Psi(\vec{r})$ for $H^{(0)}$ by forming a completely antisymmetrized spatial wavefunction from the $\phi_{m\nu}$ (as required by the FSP assumption). + This total wavefunction Wor) cau then be used along with IIO of equation (3)) to solve for the energy of the total Ixuniltonian using simple lst order perturbation theory., This total wavefunction $\Psi(\vec{r})$ can then be used along with $H^{(1)}$ of equation \ref{5.2}) ) to solve for the energy of the total Hamiltonian using simple 1st order perturbation theory. + This approach is. iu effect. a double application of perturbation theory.," This approach is, in effect, a double application of perturbation theory." + Iu the first application perturbation theory is used on the siugle-orbital Tamiltouian to handle the higher Laudau levels aud solve for the iulti-configurational. sinele-orbital energies aud wavefuuctious (ith ouly 1 iterative calculation required).," In the first application perturbation theory is used on the single-orbital Hamiltonian to handle the higher Landau levels and solve for the multi-configurational, single-orbital energies and wavefunctions (with only 1 iterative calculation required)." + Iu the second application lst order perturbation theory on the total wavefunction aud total Παπποια vields the total eucrev., In the second application 1st order perturbation theory on the total wavefunction and total Hamiltonian yields the total energy. + This same wavefunction is used το calculate oscillator streneths., This same wavefunction is used to calculate oscillator strengths. + To avoid confusion we call the first application perturbation theory of type T and the secoud application perturbation theory of type IL., To avoid confusion we call the first application perturbation theory of type I and the second application perturbation theory of type II. + This is simular to the Zexpansion method in which the whole clectrou-clectrou interaction terim is treated as a perturbation (IIvlleraas1930)., This is similar to the Z–expansion method in which the whole electron-electron interaction term is treated as a perturbation \citep{hylleras30}. +. However. the Zexpansion method is accurate only for highly ionized atoms.," However, the Z–expansion method is accurate only for highly ionized atoms." + Our Oth order is much more accurate for multi-electron atoms due to its sel£fconsistenev., Our 0th order is much more accurate for multi-electron atoms due to its self-consistency. + As a result. we achieved fast convergence (3..5 iterations for less than convereence in total cucreies) aud nunerical stability over all the electron configurations we have computedfor Z—126 atoms.," As a result, we achieved fast convergence (3–5 iterations for less than convergence in total energies) and numerical stability over all the electron configurations we have computedfor $Z=1-26$ atoms." + The method preseuted in this article is valid to oy21. covering the typical maguctic &Selds of cooling ucutron stars for the surface elements up to Fe (figure 1)).," The method presented in this article is valid to $\beta_Z > 1$, covering the typical magnetic fields of cooling neutron stars for the surface elements up to Fe (figure \ref{regime}) )." + The single-orbital Παπάοσα 2!!! in (6)) is divided iuto the Oth order and its perturbation as follows., The single-orbital Hamiltonian $h^{(0)}$ in \ref{5.4}) ) is divided into the 0th order and its perturbation as follows. +" where πυρ.1{ο) isthe transverse IEuniltonian. whose cigeuvalie and eisenfuuction are sip aud 0,,(p.0) (seeoO below)."," where $h_B(\rho, \phi)$ is the transverse Hamiltonian, whose eigenvalue and eigenfunction are $n\hbar\omega_B$ and $\Phi_{nm}(\rho,\phi)$ (see below)." +" Due to the separation of variables in5/9, eieeufunctious and eigeuvalues of the Schróddiuger equation for 0) are eiven by Ελ isa complete set of orthogonal eigeufunctious of PW "," Due to the separation of variables in$\tilde h^{(0)}$ , eigenfunctions and eigenvalues of the Schröddinger equation for $\tilde h^{(0)}$ : are given by $\chi_{nm\nu}$ is a complete set of orthogonal eigenfunctions of $\tilde h^{(0)}$ ." +egy is an energy. eigenvalue obtained by tle sinele- calculation for η=ϐ (see below)., $\epsilon_{0m\nu}$ is an energy eigenvalue obtained by the single-configurational calculation for $n=0$ (see below). +" The transverse &,,,,(9.0) is a Landau function defined as (Landau&Lifshitz 1965).. satistving the orthonormal coudition"," The transverse $\Phi_{nm}(\rho, \phi)$ is a Landau function defined as \citep{landau65}, , satisfying the orthonormal condition" +More than ninety vears ago Edcington(1918). had already realized that the period. changes of a Cepheid-like pulsator would. give information on the changes of the physics of the star's interior during its evolution.,More than ninety years ago \citet{ed18} had already realized that the period changes of a Cepheid-like pulsator would give information on the changes of the physics of the star's interior during its evolution. + The potential of measuring the rate of stellar evolution. seemed. to. be fulfilled) when Martin(1938) discovered. that the periods of the RRab stars in omega Centauri were predominantly increasing., The potential of measuring the rate of stellar evolution seemed to be fulfilled when \citet{ma38} discovered that the periods of the RRab stars in omega Centauri were predominantly increasing. + The subsequent observations of RR Lyrac stars in other globular clusters. however. failed. to demonstrate he direct connection between the detected. period. changes and horizontal-branch (1B) evolution.," The subsequent observations of RR Lyrae stars in other globular clusters, however, failed to demonstrate the direct connection between the detected period changes and horizontal-branch (HB) evolution." + The observed. large »eriod. decreases. the random and/or abrupt period changes could not be reconciled with steady evolutionary effects.," The observed large period decreases, the random and/or abrupt period changes could not be reconciled with steady evolutionary effects." + Lt ias been. however. argued that evolutionary period changes must be present in individual stars and must be observable. at least on long time-scales.," It has been, however, argued that evolutionary period changes must be present in individual stars and must be observable, at least on long time-scales." +CGlobular clusters with dozens. 1unclreds of RAR. Lyrae stars are ideal targets for such studies. the mean value of the period-change rates of a large sample of Hlt Lyrac stars of a elobular cluster must. show the trenel of their evolution.,"Globular clusters with dozens, hundreds of RR Lyrae stars are ideal targets for such studies, the mean value of the period-change rates of a large sample of RR Lyrae stars of a globular cluster must show the trend of their evolution." + Indeed. Lee(1991) proved. that the mean period-change rates observed in globular clusters were consistent with svnthetic LIB models. if the HED-tvpes of the clusters were also considered.," Indeed, \citet{le91} proved that the mean period-change rates observed in globular clusters were consistent with synthetic HB models, if the HB-types of the clusters were also considered." + The generally adopted: procedure for investigating the veriock changes of variable stars is the construction and interpretation of their ο€ diagrams., The generally adopted procedure for investigating the period changes of variable stars is the construction and interpretation of their $O-C$ diagrams. + This plots the difference between the observed. times (QO) of a particular shase. usually the maximum. or the mic-point on the ascending branch of the light curve anc the predicted ime of the same phase (€). calculated according to an accepted. ephemeris.," This plots the difference between the observed times $O$ ) of a particular phase, usually the maximum or the mid-point on the ascending branch of the light curve and the predicted time of the same phase $C$ ), calculated according to an accepted ephemeris." + However. using today's computational echniques. the ο6 data are often caleulated as the time shift between the entire observed light curve and the normal ight curve. rather than from one specific phase.," However, using today's computational techniques, the $O-C$ data are often calculated as the time shift between the entire observed light curve and the normal light curve, rather than from one specific phase." + This method gives more stable results., This method gives more stable results. + Hf the ο—€ diagram is a linear, If the $O-C$ diagram is a linear +"Out-of-pupil optics could have a strong impact on the performances of any differential technique adopted in high-contrast imaging due to Fresnel propagation, as described by Maroisetal.(2006).","Out-of-pupil optics could have a strong impact on the performances of any differential technique adopted in high-contrast imaging due to Fresnel propagation, as described by \citet{Ma06}." +. An optic that is not conjugated to a pupil plane will modify the light distribution in a chromatic way because at this location the beam intensity distribution depends on wavelength through diffraction effects., An optic that is not conjugated to a pupil plane will modify the light distribution in a chromatic way because at this location the beam intensity distribution depends on wavelength through diffraction effects. +" The closer the optic is to a focal plane, the larger this chromaticity."," The closer the optic is to a focal plane, the larger this chromaticity." +" Even more severe is the fact that this chromaticity is no longer smooth, but cyclic along the spectrum, when the optic is conjugated to a height that is several times the Talbot length defined as where 4 is the light wavelength and A the period of a single sinusoidal component of the wavefront across the pupil."," Even more severe is the fact that this chromaticity is no longer smooth, but cyclic along the spectrum, when the optic is conjugated to a height that is several times the Talbot length defined as where $\lambda$ is the light wavelength and $\Lambda$ the period of a single sinusoidal component of the wavefront across the pupil." +" For an aberration with a given period, the pupil complex amplitude presenting the electromagnetic field changes from a pure wavefront error to a pure amplitude error over a quarter of the Talbot length."," For an aberration with a given period, the pupil complex amplitude presenting the electromagnetic field changes from a pure wavefront error to a pure amplitude error over a quarter of the Talbot length." +" Since the Talbot length is different for different periods, a decorrelation occurs that depends on angular separation."," Since the Talbot length is different for different periods, a decorrelation occurs that depends on angular separation." +" The farther an optic is from the pupil plane (in multiples of Talbot length), the more the decorrelation along spectral domain will be, and speckle correlation will be In the case of SPHERE, the Talbot effect was expected to be strong for those optical components located before the lenslet array, such as the entrance window, the ADC, the derotator and the coronagraphic mask (Yaitskovaetal.2010)."," The farther an optic is from the pupil plane (in multiples of Talbot length), the more the decorrelation along spectral domain will be, and speckle correlation will be In the case of SPHERE, the Talbot effect was expected to be strong for those optical components located before the lenslet array, such as the entrance window, the ADC, the derotator and the coronagraphic mask \citep{Yai10}." +". To evaluate the impact of the Fresnel propagation, we cannot use the CAOS package that is based on the Fraunhofer propagation."," To evaluate the impact of the Fresnel propagation, we cannot use the CAOS package that is based on the Fraunhofer propagation." + We then exploited the PROPER code to create new PSFs that are then used as input for the CSP code., We then exploited the PROPER code \citep{Kr07} to create new PSFs that are then used as input for the CSP code. + In Table 2 we list the parameters used to calculate the Fresnel propagation in all the simulations., In Table \ref{tabfresnelparameter} we list the parameters used to calculate the Fresnel propagation in all the simulations. + We report the values of the conjugated distance and the wavefront error (WFE) rms for all the considered optical surfaces., We report the values of the conjugated distance and the wavefront error (WFE) rms for all the considered optical surfaces. +" To save computing time, we performed all these simulations without considering the effects of the atmosphere (using only 1 atmospheric phase screen)."," To save computing time, we performed all these simulations without considering the effects of the atmosphere (using only 1 atmospheric phase screen)." +" Of course, this is not realistic, because it yields contrasts that are to optimistics."," Of course, this is not realistic, because it yields contrasts that are to optimistics." +" However, the comparison is still meaningful for evaluating the impact of Fresnel propagation From comparison of the plots in Figs."," However, the comparison is still meaningful for evaluating the impact of Fresnel propagation From comparison of the plots in Figs." +" 8 and 9 resulting from simulations that do not include and that do include the Fresnel propagation effects, respectively, one can see that the differences between the two cases are very small, as confirmed by the data reported in Table 3 where the values of the contrast at different separations for simulations performed not considering (second column) and considering (third column)"," \ref{nofresnel} and \ref{fresnel} resulting from simulations that do not include and that do include the Fresnel propagation effects, respectively, one can see that the differences between the two cases are very small, as confirmed by the data reported in Table \ref{tabdatafresnel} where the values of the contrast at different separations for simulations performed not considering (second column) and considering (third column)" +"bars, Rote~0.027 Gyr-!.","bars, $R_{\rm mm}^{\rm blue} \sim 0.027$ $^{-1}$." +" To the contrary, the major mergerrate, denoted Rb, decreases by a factor of three from z=0.8 to z=0.5, as noted by?."," To the contrary, the major mergerrate, denoted $R_{\rm MM}^{\rm blue}$, decreases by a factor of three from $z = 0.8$ to $z = 0.5$, as noted by." +". These trends suggest that the stability or increase with cosmic time of the minor merger rate found in the previous section is a consequence of the evolution in the fraction of bright galaxies that are red: as time goes by, the red fraction increases??)."," These trends suggest that the stability or increase with cosmic time of the minor merger rate found in the previous section is a consequence of the evolution in the fraction of bright galaxies that are red: as time goes by, the red fraction increases." +". Because the minor merger rate of red galaxies is a factor of ~2.5 higher than the one of blue galaxies, and both are roughly constant, the increase in the red fraction implies an increase in the global (red+blue) minor merger rate."," Because the minor merger rate of red galaxies is a factor of $\sim2.5$ higher than the one of blue galaxies, and both are roughly constant, the increase in the red fraction implies an increase in the global (red+blue) minor merger rate." +" This effect is also present in the major merger rate, but in this case Ro decreases with cosmic time, and the increase in the red fraction is only a mild evolution, as found by?."," This effect is also present in the major merger rate, but in this case $R_{\rm MM}^{\rm blue}$ decreases with cosmic time, and the increase in the red fraction is only a mild evolution, as found by." +". The volumetric merger rate (i.e., the number of mergers per unit volume and time) is a complementary measure to the merger rate estimated in the previous sections."," The volumetric merger rate (i.e., the number of mergers per unit volume and time) is a complementary measure to the merger rate estimated in the previous sections." +" To obtain the volumetric merger rate, denoted ®, we multiply the merger rate by the number density of all/red/blue galaxies with ΜΕx—20 in VVDS-Deep at each redshift(?)."," To obtain the volumetric merger rate, denoted $\Re$, we multiply the merger rate by the number density of all/red/blue galaxies with $M_B^{\rm e} \leq -20$ in VVDS-Deep at each redshift." +. We summarize the values of ® in Table 9.., We summarize the values of $\Re$ in Table \ref{mrvtab}. + All trends are similar to those found in the previous section., All trends are similar to those found in the previous section. +" Interestingly, we find that Rd.3qbhe~3.5x10?Mpc? Gyr!."," Interestingly, we find that $\Re_{\rm mm}^{\rm red} \sim \Re_{\rm mm}^{\rm blue} \sim 3.5 \times 10^{-5}\ {\rm Mpc}^{-3}\ {\rm Gyr}^{-1}$ ." +" The merger rate of red galaxies is higher by a factor of ~ 2.5 than that of the blue ones, but the number density of the latter is higher than of the former, hence making the volumetric merger rates comparable."," The merger rate of red galaxies is higher by a factor of $\sim$ 2.5 than that of the blue ones, but the number density of the latter is higher than of the former, hence making the volumetric merger rates comparable." +" In this section we estimate the evolution of the minor merger fraction and rate with redshift, and discuss the contribution of minor mergers to the evolution of bright galaxies since z~1, comparing it to the contribution of major mergers."," In this section we estimate the evolution of the minor merger fraction and rate with redshift, and discuss the contribution of minor mergers to the evolution of bright galaxies since $z \sim 1$, comparing it to the contribution of major mergers." +" The evolution of the merger faction with redshift up to is well parametrized by a power-law???),, Our results alone suggest that the merger fractionevolves faster for higher u, with πι=5.6 for equal luminosity companions (u= 1), m=2.4 for major companions with u> 1/4, and m=0.8 for major + minor companions (u> 1/10)."," The evolution of the merger faction with redshift up to $z \sim 1.5$ is well parametrized by a power-law, Our results alone suggest that the merger fractionevolves faster for higher $\mu$ , with $m=5.6$ for equal luminosity companions $\mu = 1$ ), $m = 2.4$ for major companions with $\mu \geq 1/4$ , and $m = 0.8$ for major + minor companions $\mu \geq 1/10$ )." + This mild evolution in the total (major + minor) merger fraction is also suggested by the morphological studies of and?., This mild evolution in the total (major + minor) merger fraction is also suggested by the morphological studies of and. +". To better constrain the evolution with redshift of the minor merger fraction, a local reference is important."," To better constrain the evolution with redshift of the minor merger fraction, a local reference is important." +" estimate that the minor merger fraction is similar to the major one (finjw~1, µ= 1/3) in Galaxy (?);; the latter is based on the visual classificationof Sloan Digital Sky Survey (SDSS?,, ?))galaxies by internet users."," estimate that the minor merger fraction is similar to the major one $f_{m/M} \sim 1$, $\mu \gtrsim 1/3$ ) in Galaxy ; the latter is based on the visual classificationof Sloan Digital Sky Survey , )galaxies by internet users." +" However, their sample is incomplete"," However, their sample is incomplete" +aba,t_0. +" Assuming.toot~LOM Lz. fy~10""m sS and taking. the other parameters the same as those in 3.1. we determine the luminosity of TeV photons produced by SSC as about Lec(TeV)~107eres"," Assuming$\nu_{\rm seed}\sim 10^{14}$ Hz, $t_0\sim 10^{9}$ s, and taking the other parameters the same as those in 3.1, we determine the luminosity of TeV photons produced by SSC as about $L_{\rm SSC}({\rm TeV})\sim 10^{30}\ {\rm +erg\ s^{-1}}$." + The present. TeV telescopes have detected: excess TeV whotons from three 5-rav. pulsars. namely. Crab. Vela aid SR. 1706-44. with a possible detection from Geminga (see reviews by Aharonian 1909: Weekes 2004).," The present TeV telescopes have detected excess TeV photons from three $\gamma$ -ray pulsars, namely, Crab, Vela and PSR 1706-44, with a possible detection from Geminga (see reviews by Aharonian 1999; Weekes 2004)." + Their >-ray emissions in the GeV band are dominated. by contributions rom the pulsar magnetosphere. as predicted in the outer gap model 2.1).," Their $\gamma$ -ray emissions in the GeV band are dominated by contributions from the pulsar magnetosphere, as predicted in the outer gap model 2.1)." + Llowever. no pulsed TeV photons are detected w the present telescopes (Ixifune et al.," However, no pulsed TeV photons are detected by the present telescopes (Kifune et al." + 1995: Yoshikoshi et al., 1995; Yoshikoshi et al. + 1997: Aharonian et al., 1997; Aharonian et al. + 1999: Lessard ct al., 1999; Lessard et al. + 2000). hus implving that the TeV photons maink come from heir wind nebulae.," 2000), thus implying that the TeV photons mainly come from their wind nebulae." + In this section. we check our one-zone model by estimating the TeV signals from these four sources.," In this section, we check our one-zone model by estimating the TeV signals from these four sources." + Because of the simplified nature of these models. we do not expect to reproduce the detailed. properties of the nebula.," Because of the simplified nature of these models, we do not expect to reproduce the detailed properties of the nebula." + Alore detailed: discussions for individual sources have been acelressecl by other authors (e.g. de Jager Harding 1992: Atovan Aharonian 1996: Aharonian. Atovan Ixifune 1991).," More detailed discussions for individual sources have been addressed by other authors (e.g. de Jager Harding 1992; Atoyan Aharonian 1996; Aharonian, Atoyan Kifune 1997)." + The Crab nebula is the first. pulsar wind. nebula with TeV signal detection. (Weekes ot al., The Crab nebula is the first pulsar wind nebula with TeV signal detection (Weekes et al. + 19590). and it. has oen extensively stucied by many LeV detectors. since. or example. Milagro (Atkins et al.," 1989), and it has been extensively studied by many TeV detectors since, for example, Milagro (Atkins et al." + 2008). LECIRA ( Aharonianet al.," 2003), HEGRA (Aharonian et al." + 2000). Whipple (Lessardet al.," 2000), Whipple (Lessard et al." + 2000). Tibet (Amenomort et al.," 2000), Tibet (Amenomori et al." + 1999). and CANCGABOO (Lamimori et al.," 1999), and CANGAROO (Tamimori et al." + L998)., 1998). + Detailed: spectral ancl ux information have oen obtained., Detailed spectral and flux information have been obtained. + Lhe Crab TeV spectrum can be fitted in he energv range above 1 ‘TeV by a simple power law with a photon index 2.5+0.2. and shows no cut-oll up to 100 TeV. The TeV luminosity is estimated. to be ~10ergs," The Crab TeV spectrum can be fitted in the energy range above 1 TeV by a simple power law with a photon index $\sim~2.5\pm 0.2$, and shows no cut-off up to 100 TeV. The TeV luminosity is estimated to be $\sim~10^{34}\ {\rm erg\ s^{-1}}$." + Atovan Aharonian (1996) discussed the radiation mechanism ofthe Crab nebula based on the model of Ixennel Coroniti (1984). and fitted the observed data from radio to GeV by svnchrotron radiation. and the TeV band through ICS) processes.," Atoyan Aharonian (1996) discussed the radiation mechanism of the Crab nebula based on the model of Kennel Coroniti (1984), and fitted the observed data from radio to GeV by synchrotron radiation, and the TeV band through ICS processes." + With a complicated spatial structure. the Crab nebula probably cannot be fully described. in. our one-zone model.," With a complicated spatial structure, the Crab nebula probably cannot be fully described in our one-zone model." + What we present in the following should therefore be treated as a rough check., What we present in the following should therefore be treated as a rough check. + The magnetic field in the Crab. nebula is relatively high at B~107 € (Lester et al., The magnetic field in the Crab nebula is relatively high at $B\sim 10^{-3}-10^{-4}$ G (Hester et al. + 1996). which significantly reduces the relative contributions Lrom the ICS processes.," 1996), which significantly reduces the relative contributions from the ICS processes." + With a ratio woam/we107. the SSC contribution to. αν photons should. be the dominant component.," With a ratio $w_{\rm CMB}/w_B\leq 10^{-4}$, the SSC contribution to TeV photons should be the dominant component." +" Assuming B~10.3 C. we derive the cooling Lorentz [actor of electrons to be =,~—10"". ancl the characteristic timescale fg~3107"" s Phe implied cooling energy 427*1l Τον which is only of the predicted SSC) cut-oll."," Assuming $B\sim 10^{-4}$ G, we derive the cooling Lorentz factor of electrons to be $\gamma_c\sim 10^6$, and the characteristic timescale $t_0\sim 3\times 10^{10}$ s. The implied cooling energy $E_c^{\rm SSC}\sim 1$ TeV, which is only of the predicted SSC cut-off." + Pherefore. above 1 TeV. he photon index of the LCOS spectrum is given by DP—(p|2)/2. where p2.," Therefore, above 1 TeV, the photon index of the ICS spectrum is given by $\Gamma=(p+2)/2$, where $p>2$ ." + On the other hand. the observed TeV spectrum of the Crab nebula implies that p.3.," On the other hand, the observed TeV spectrum of the Crab nebula implies that $p\sim 3$ ." + However. from. the observed:svnchrotron X-ray spectrum. of the Crab nebula (Willingale et al.," However, from the observedsynchrotron X-ray spectrum of the Crab nebula (Willingale et al." + 2001). P=(p|2)/2—2.1. which gives po2.2.," 2001), $\Gamma=(p+2)/2\sim 2.1$, which gives $p\sim 2.2$." +" Since 52004;75s010"" in the Crab nebula. and above terimic?~L TeV. the SSC spectrum can be steeper than (p|2)/2. hence agreeing with the observations."," Since $\gamma_{crit}\sim \gamma_c\sim 10^6$ in the Crab nebula, and above $\gamma_{crit} m_ec^2\sim 1$ TeV, the SSC spectrum can be steeper than $(p+2)/2$, hence agreeing with the observations." + The discrepancy may also result from the spatial variations of the electron. spectrum and. magnetic field in the Crab nebula considered. by other models (ce Jager Harding. 1992: Atovan Aharonian 1996: Sollerman et al., The discrepancy may also result from the spatial variations of the electron spectrum and magnetic field in the Crab nebula considered by other models (de Jager Harding 1992; Atoyan Aharonian 1996; Sollerman et al. + 2000). which ave not considered in our one-zone model.," 2000), which are not considered in our one-zone model." + Despite the very. large size of the Crab. nebula. the JeV photons come from a central region 1 pe (Atkins et al.," Despite the very large size of the Crab nebula, the TeV photons come from a central region $\sim 1$ pc (Atkins et al." + 2003)., 2003). + The termination radius of the nebula is about H.—10'4T em.," The termination radius of the nebula is about $R_s=4\times +10^{17}$ cm." +" In our one-zone model. we can take Goep~OS.-p=23. and 5,4~10"". to estimate the svnehrotron Luminosity and energy density."," In our one-zone model, we can take $\epsilon_e\sim\epsilon_B\sim 0.5, \ p=2.2$, and $\gamma_w\sim +10^7$, to estimate the synchrotron luminosity and energy density." + Atv ~Lott Ην. the Duminosity Loa~107Eergs and tyορ~0.03.," At $\nu\sim 10^{14}$ Hz, the luminosity $L_{\rm syn}\sim 10^{36}\ +{\rm erg\ s^{-1}}$ and $w_{\rm syn}/w_B\sim 0.03$." +2 Therefore the SSC contribution to TeV luminosity is Lisc~3.⋅LO3dergs1 . which. is. comparable to the observed value.," Therefore the SSC contribution to TeV luminosity is $L_{\rm SSC}\sim 3\times +10^{34}\ {\rm erg\ s^{-1}}$ , which is comparable to the observed value." + Yoshikoshi et al. (, Yoshikoshi et al. ( +1997) used the 3.8 m imaging C'erenkov telescope near Woomera. South Australia to detect TeV photons of the Vela pulsar region.,"1997) used the 3.8 m imaging Cerenkov telescope near Woomera, South Australia to detect TeV photons of the Vela pulsar region." +" The detected. TeV + -rav Hux is 2.910""photonem7s|. corresponding to a luminosity of 107ergs.I assuming the distance. of Vela to be 300 pe."," The detected TeV $\gamma$ -ray flux is $2.9\times +10^{-12}{\rm photon\ cm^{-2}\ s^{-1}}$, corresponding to a luminosity of $\times 10^{32}\ {\rm erg\ s^{-1}}$ assuming the distance of Vela to be 300 pc." + Similar to Vela. PSR. 1706-44 has also been detected in the TeV energy range (Ixifune et al.," Similar to Vela, PSR 1706-44 has also been detected in the TeV energy range (Kifune et al." + 1995: Chadwick et al., 1995; Chadwick et al. + 1998)., 1998). + The Ilux above a threshold of 1 Tey is —10++photonem?s.+., The flux above a threshold of 1 TeV is $\sim 10^{-11}\ {\rm photon\ cm^{-2}\ s^{-1}}$. + However. recent observations of HESS reported non-detection of PSR 1706-4. with an upper limit flux ~L3«10“ereem?s (Aharonian et al.," However, recent observations of HESS reported non-detection of PSR 1706-44, with an upper limit flux $\sim 1.3\times 10^{-12}{\rm erg\ cm^{-2}\ +s^{-1}}$ (Aharonian et al." + 2005)., 2005). + The reason for the observational discrepancy is unknown. possibly the “TeV emission. could be variable on a time scale of vears. e.g. in the different mediunr environments.," The reason for the observational discrepancy is unknown, possibly the TeV emission could be variable on a time scale of years, e.g. in the different medium environments." + Assuming a distance of 1.5 kpe. the TeV luminosity is lower than 4.1077ergs," Assuming a distance of 1.5 kpc, the TeV luminosity is lower than $4\times 10^{32}\ {\rm erg\ s^{-1}}$." + TeV spectral information has been obtained for neither of these two SOULCCS., TeV spectral information has been obtained for neither of these two sources. +" Vela and PSR. 1706-44 have nearly the same spin down power of a few times LO""Eergs 1.", Vela and PSR 1706-44 have nearly the same spin down power of a few times $10^{36}\ {\rm erg\ s^{-1}}$ . + This suggests that their nebulae have similar propertics moving in the interstellar medium., This suggests that their nebulae have similar properties moving in the interstellar medium. + Assumine¢.~0.5andcg 0.001. we estimate the magnetic field D in the two pulsar wind nebulae.," Assuming $\epsilon_e \sim 0.5$and$\epsilon_B \sim 0.001$ , we estimate the magnetic field $B$ in the two pulsar wind nebulae." + Xecording to the Chandra observation of Vela (Pavlov et al., According to the Chandra observation of Vela (Pavlov et al. + 2001).," 2001)," +these prescriptions are ad hoc. they may still capture oeuportaut aspects of the plivsies iu real astrophysical flows aud provide useful coustraimts by direct comparison with observations.,"these prescriptions are ad hoc, they may still capture important aspects of the physics in real astrophysical flows and provide useful constraints by direct comparison with observations." + With this motivation. we lave! computed the accretion isk spectra preclicted by several differeut stress prescriptions using the most colplete spectral code currently available.," With this motivation, we have computed the accretion disk spectra predicted by several different stress prescriptions using the most complete spectral code currently available." + This inchides both non-LTE ion populations. with radiative traiister iu the disk aud full eoncral relativistic rav tracing to propagate this flux to the observer.," This includes both non-LTE ion populations, with radiative transfer in the disk and full general relativistic ray tracing to propagate this flux to the observer." + These models eive results which are very close to the observed behavior of the spectra of black hole binaries iu the therial dominant (ligh/soft) state from L/Lpgg;~0.06»(0.6., These models give results which are very close to the observed behavior of the spectra of black hole binaries in the thermal dominant (high/soft) state from $L/L_{Edd}\sim 0.06\to 0.6$. +" The majority of this data comes from provortional counters. anc eoncrally shows that the temperature changes slightly more rapidly with ΕΠ fhan expected for a disk of constant iuner radius. aud costant color-teniperature correction (e.g. the compilation of DCOL and Ὀ(νο,"," The majority of this data comes from proportional counters, and generally shows that the temperature changes slightly more rapidly with luminosity than expected for a disk of constant inner radius, and constant color-temperature correction (e.g. the compilation of DG04 and DGK07)." + This is as predicted by all the Heh surface density disk models presented here. auc should hold generally for all stress prescriptions where less than 10 per cent of the energv is dissipated above the effective photosphere.," This is as predicted by all the high surface density disk models presented here, and should hold generally for all stress prescriptions where less than 10 per cent of the energy is dissipated above the effective photosphere." + The surface lavers then simply act like a passive atmosphere. with properties set bv the effective temperature.," The surface layers then simply act like a passive atmosphere, with properties set by the effective temperature." + The increase in temperature with luminosity gives rise to a sanall inerease m color-teniperature correction and the resulting spectra are remarkably similar respective of the detailed form of the stress., The increase in temperature with luminosity gives rise to a small increase in color-temperature correction and the resulting spectra are remarkably similar irrespective of the detailed form of the stress. + Such inodel fitting can also place interesting constraints on the angular momentum transport., Such model fitting can also place interesting constraints on the angular momentum transport. + The classic alpha disk with a=00.1 becomes effectively optically thin at the highest huuimosities., The classic alpha disk with $\alpha=0.1$ becomes effectively optically thin at the highest luminosities. + The resulting increase iu color temperature with hnuuiuositv is much more rapid than observed from existing data. ruling out such stress prescriptions.," The resulting increase in color temperature with luminosity is much more rapid than observed from existing data, ruling out such stress prescriptions." + The standard alpha disk i also uustable to the thermalviscous radiation pressure instability which predicts hut cycle behavior which is uot observed. again showing that the disk is deuser than such models predict.," The standard alpha disk is also unstable to the thermal-viscous radiation pressure instability which predicts limit cycle behavior which is not observed, again showing that the disk is denser than such models predict." + However. even these models eive the same spectra at low luninosities (L.XO12baa). where the surface density is sufficiently large for the disk to remain very effectively optically thick.," However, even these models give the same spectra at low luminosities $L \lesssim 0.1 L_{\rm Edd}$ ), where the surface density is sufficiently large for the disk to remain very effectively optically thick." + Therefore. where we lave a clear view of the disk. unaffected by a anoderately optically thick disk wiud. the disk spectra should provide a relatively robust estimator of the disk inner radius. aud plausibly. the spin of the black hole.," Therefore, where we have a clear view of the disk, unaffected by a moderately optically thick disk wind, the disk spectra should provide a relatively robust estimator of the disk inner radius, and plausibly, the spin of the black hole." +The faint enmuüssiou of transicuthy accreting low-mass bbiuaries in quiescence (qLAINBs) found iu globular clusters (CC's) ds often used to measure the τας of neutron stars (NSs).,The faint emission of transiently accreting low-mass binaries in quiescence (qLMXBs) found in globular clusters (GCs) is often used to measure the radii of neutron stars (NSs). + Such measurements. free of large uncertainties caused by uukuowu distances aud mospleric composition. can provide useful constraints on the nuclear dense matter equation of state (EoS) relating pressure aud density when matter has a density above 2.35«105e7. such as can occur iu atomic wielei. and iu the interiorsc1. of NSs (?)..," Such measurements, free of large uncertainties caused by unknown distances and atmospheric composition, can provide useful constraints on the nuclear dense matter equation of state (EoS) relating pressure and density when matter has a density above $2.35\tee{14}\cgsdensity$, such as can occur in atomic nuclei, and in the interiors of NSs \citep{lattimer04}." + Obtaining precise constraiuts on the deuse matter Eos is the observational uotivatiou for NS radii micasurements. requiring at least ac5 accuracy to be useful (?7)..," Obtaining precise constraints on the dense matter EoS is the observational motivation for NS radii measurements, requiring at least a $\sim5\%$ accuracy to be useful \citep{lattimer04}." + The expected overabundance of LAUINBs in GCs (7) las motivated observations with the current gcucration of telescopes., The expected overabundance of LMXBs in GCs \citep{hut92} has motivated observations with the current generation of telescopes. + It was then empirically shown that the NS ünaries population in GC depends on the interaction vate of the cluster (?777)..," It was then empirically shown that the NS binaries population in GC depends on the interaction rate of the cluster \citep{gendre03b,heinke03b,pooley03}." + Several qLMXDs were discovered close to the cores of the GCs (forexample. 7)..," Several qLMXBs were discovered close to the cores of the GCs \citep[for example,][]{rutledge02b}." + In inost. cases. they were spectrally identified based ou their spectra consistent with a NS atinosphere at the distance of their lost cluster.," In most cases, they were spectrally identified based on their spectra consistent with a NS atmosphere at the distance of their host cluster." + The known distances to he clusters aud the known values of the absorption led o precise NS radii measurements (~5—20%uncor-ainty. 72773.," The known distances to the clusters and the known values of the absorption led to precise NS radii measurements \citep[$\sim 5-20\%$ +uncertainty, ][]{rutledge02b, gendre03b, heinke03c}." + Tlowever. oulv a haudful of coufirmed CC qLMXDs are known. aud finding more of those objects recessitates careful oobservatious of GCs.," However, only a handful of confirmed GC qLMXBs are known, and finding more of those objects necessitates careful observations of GCs." + Finally. the X-ray uuuositv qLMXDs is not expected to show variability Oll NORER iuescales (2°).. but loue-term flux variation have been observed before (2)..," Finally, the X-ray luminosity qLMXBs is not expected to show variability on $\sim$ years timescales \citep{brown98,ushomirsky01}, but long-term flux variation have been observed before \citep{rutledge02a}." +. A amore complete introduction about qLAINBs and an exhaustive list of GC qLAINBs can be found elsewhere (2.G09hereatter).., A more complete introduction about qLMXBs and an exhaustive list of GC qLMXBs can be found elsewhere \citep[G09 hereafter]{guillot09a}. + Regarding the observations of GC's. the two telescopes aud: aare colmplementary.," Regarding the observations of GCs, the two telescopes and are complementary." +" On the one hand. the typical unnmnositv of qLAINBs (Ex~107.10eres 1) and he typical distances to GCs require the effective area of to collect high signal-to-noise ratio (S/N) data for spectral analyses. within modest integration times,"," On the one hand, the typical luminosity of qLMXBs $L_{\rm X}\sim10^{32}-10^{33} \cgslum$ ) and the typical distances to GCs require the effective area of to collect high signal-to-noise ratio (S/N) data for spectral analyses, within modest integration times." + Ou he other haud. Chendrss angular resolution. pcriits spatial resolution of adjaceut sources im cores of GC. which," On the other hand, s angular resolution permits spatial resolution of adjacent sources in cores of GC, which" +water line emission in the HIFI beam.,water line emission in the HIFI beam. + Follow-up observations of the higher-exeitation lines of CO and wwith HIFI will help us constrain more accurately the physical conditions of each velocity component (density. temperature) and more generally in the shock.," Follow-up observations of the higher-excitation lines of CO and with HIFI will help us constrain more accurately the physical conditions of each velocity component (density, temperature) and more generally in the shock." + ΠΙΕΙ has been designed and built. by a consortium of institutes and university departments from across Europe. Canada and the United States under the leadership of SRON Netherlands Institute for Space Research. Groningen. The Netherlands and with major contributions from Germany. France and the US.," HIFI has been designed and built by a consortium of institutes and university departments from across Europe, Canada and the United States under the leadership of SRON Netherlands Institute for Space Research, Groningen, The Netherlands and with major contributions from Germany, France and the US." + Consortium members are: Canada: CSA. U.Waterloo: France: CESR. LAB. LERMA. IRAM: Germany: KOSMA. MPIfR. MPS: Ireland. NUI Maynooth: Italy: ASI. IFSI-INAF. Osservatorio Astrofisico di Arcetri- INAF: Netherlands: SRON. TUD: Poland: CAMK. CBK: Spain: Observatorio Astronómmico Nacional (IGN). Centro de Astrobiologiaa (CSIC-INTA).," Consortium members are: Canada: CSA, U.Waterloo; France: CESR, LAB, LERMA, IRAM; Germany: KOSMA, MPIfR, MPS; Ireland, NUI Maynooth; Italy: ASI, IFSI-INAF, Osservatorio Astrofisico di Arcetri- INAF; Netherlands: SRON, TUD; Poland: CAMK, CBK; Spain: Observatorio Astronómmico Nacional (IGN), Centro de a (CSIC-INTA)." + Sweden: Chalmers University of Technology - MC2. RSS GARD: Onsala Space Observatory: Swedish National Space Board. Stockholm University - Stockholm Observatory: Switzerland: ΕΤΗ Zurich. FHNW: USA: Caltech. JPL. NHSC.," Sweden: Chalmers University of Technology - MC2, RSS GARD; Onsala Space Observatory; Swedish National Space Board, Stockholm University - Stockholm Observatory; Switzerland: ETH Zurich, FHNW; USA: Caltech, JPL, NHSC." + HIPE is a joint development by the Science Ground Segment Consortium. consisting of ESA. the NASA Science Center. and the HIFI. PACS and SPIRE consortia.," HIPE is a joint development by the Science Ground Segment Consortium, consisting of ESA, the NASA Science Center, and the HIFI, PACS and SPIRE consortia." +a ‘strange’ galaxy with two counter-rotating dises (Bertola et al.,a `strange' galaxy with two counter-rotating discs (Bertola et al. + 1996)., 1996). +" We find some hint that higher 7,45,/7. ratios are found in smaller galaxies and vice versa.", We find some hint that higher $r_{\sigma\rm{-drop}}/r_{\rm{c}}$ ratios are found in smaller galaxies and vice versa. + That would indicate that rícapgalaxy. which ts confirmed by the plot in the right panel of Fig. 2..," That would indicate that $r_{\sigma\rm{-drop}}$, which is confirmed by the plot in the right panel of Fig. \ref{sigmasizes}." + Because we ignored bars in our sample selection. procedure (Sect., Because we ignored bars in our sample selection procedure (Sect. + 2). We can as a first step investigate the presence and properties of the bars in our c-drop and control sample galaxies.," 2), we can as a first step investigate the presence and properties of the bars in our $\sigma$ -drop and control sample galaxies." + The interest in this Hes primarily in the fact. that bars can facilitate the inflow of cool gas into the inner regions of galaxies by removing angular momentum from the gas in the outer parts., The interest in this lies primarily in the fact that bars can facilitate the inflow of cool gas into the inner regions of galaxies by removing angular momentum from the gas in the outer parts. + If o-drops are related to star formation from dynamically cold gaseous material close to the nucleus. it is reasonable to conjecture that bars may play an important role in the process.," If $\sigma$ -drops are related to star formation from dynamically cold gaseous material close to the nucleus, it is reasonable to conjecture that bars may play an important role in the process." + In Table 2 and Table 3 we therefore summarise the bar classification of the galaxies., In Table \ref{sigmaproperties} and Table \ref{controlproperties} we therefore summarise the bar classification of the galaxies. + This is based primarily on the RC3 catalogue (de Vaucouleurs et al., This is based primarily on the RC3 catalogue (de Vaucouleurs et al. + 1991). where. as Is conventional. SA galaxies are those without a bar. SAB (also known as SX) are weakly or moderately barred. and SB galaxies have a prominent bar in the optical.," 1991), where, as is conventional, SA galaxies are those without a bar, SAB (also known as SX) are weakly or moderately barred, and SB galaxies have a prominent bar in the optical." + Although the RC3 is based on optical imaging and is therefore not optimal in recognising all bars. it is generally reliable in reporting well-defined bars.," Although the RC3 is based on optical imaging and is therefore not optimal in recognising all bars, it is generally reliable in reporting well-defined bars." + We did check the literature and images of individual galaxies though. and refined the classification in the following cases.," We did check the literature and images of individual galaxies though, and refined the classification in the following cases." + Some galaxies are catalogued as non-barred in the RC3 but do in fact have weak and/or small bars., Some galaxies are catalogued as non-barred in the RC3 but do in fact have weak and/or small bars. + The first case is NGC 1068. part of our c-drop sample. where a bar has been discovered from near-IR imaging (e.g.. Scovilleetal. 1988)).," The first case is NGC 1068, part of our $\sigma$ -drop sample, where a bar has been discovered from near-IR imaging (e.g., \cite{SC88}) )." + Other galaxies with a bar discovered on the basis of modern imaging are the control sample galaxies NGC 3169. (Laurikainen&Salo 2002)). NGC 3593 and NGC 3675 (Laurikainenetal. 2004)).. and NGC 2460. NGC 7331. and NGC 7742 (Laineetal. 2002)).," Other galaxies with a bar discovered on the basis of modern imaging are the control sample galaxies NGC 3169 \cite{LA02}) ), NGC 3593 and NGC 3675 \cite{LA04}) ), and NGC 2460, NGC 7331, and NGC 7742 \cite{LAI02}) )." + Fitting ellipses to à 2MASS image we also found strong evidence for a weak bar in NGC 2985., Fitting ellipses to a $2MASS$ image we also found strong evidence for a weak bar in NGC 2985. + The bar fractions 1n our c-drop and control samples can thus be seen in Table 6.. where we use Poisson statistics to give an estimate of the error in the numbers (e=VFO-FIN). where f is the quantity that is measured and N is the sample size in which the quantity is searched).," The bar fractions in our $\sigma$ -drop and control samples can thus be seen in Table \ref{bars}, where we use Poisson statistics to give an estimate of the error in the numbers $\epsilon=\sqrt{f(1-f/N)}$, where $f$ is the quantity that is measured and $N$ is the sample size in which the quantity is searched)." + The conclusion from the numbers presented in Table 6 is that there are no statistically significant differences between the two samples., The conclusion from the numbers presented in Table \ref{bars} is that there are no statistically significant differences between the two samples. + What is evident from this test. however. is that a bar is not needed to cause a c-drop. and that. conversely. the presence of a bar does not automatically. or even preferentially. lead to à dip m the central stellar velocity dispersion.," What is evident from this test, however, is that a bar is not needed to cause a $\sigma$ -drop, and that, conversely, the presence of a bar does not automatically, or even preferentially, lead to a dip in the central stellar velocity dispersion." + As for inner bars. we have found six of them in c-drop galaxies (NGC 1068. NGC 1097. NGC 1808. NGC 2460. NGC 4303. and NGC 4725) and three in galaxies of the control sample (namely NGC 3169. NGC 3368. and NGC 7742).," As for inner bars, we have found six of them in $\sigma$ -drop galaxies (NGC 1068, NGC 1097, NGC 1808, NGC 2460, NGC 4303, and NGC 4725) and three in galaxies of the control sample (namely NGC 3169, NGC 3368, and NGC 7742)." + All of these inner bars have been reported independently by Laine et al. (, All of these inner bars have been reported independently by Laine et al. ( +2002) and/or Erwin (2004).,2002) and/or Erwin (2004). + The difference between inner bar fraction in the c-drop and control samples is not very significant. but points m the direction of a link between inner bars and c-drops.," The difference between inner bar fraction in the $\sigma$ -drop and control samples is not very significant, but points in the direction of a link between inner bars and $\sigma$ -drops." + It is worthy to note here that. except in the case of NGC 3593. cr-drops have a much smaller scale than outer bars.," It is worthy to note here that, except in the case of NGC 3593, $\sigma$ -drops have a much smaller scale than outer bars." + In the other hand. c-drops occur on a size scale similar to that of inner bars (our nine inner bars have a mean size of ppc. compared to the mean radius of the c-drops in the sample. which is ppc).," In the other hand, $\sigma$ -drops occur on a size scale similar to that of inner bars (our nine inner bars have a mean size of pc, compared to the mean radius of the $\sigma$ -drops in the sample, which is pc)." +"N-bocly systems. so we cannot expect simulations of the ""ame object run with dillerent. resolution. with cillerent codes. or with dillerent. integration parameters to be very similar at the final time. (","N-body systems, so we cannot expect simulations of the `same' object run with different resolution, with different codes, or with different integration parameters to be very similar at the final time. (" +See for example the various simulations from identical initial conditions in the Santa Barbara Cluster Comparison Project’ (Frenk et al.,See for example the various simulations from identical initial conditions in the Santa Barbara Cluster Comparison Project (Frenk et al. + 1999) This is because in a chaotic N-bocky system any small »erturbation to the trajectory is amplified exponentially by subsequent evolution., 1999) This is because in a chaotic N-body system any small perturbation to the trajectory is amplified exponentially by subsequent evolution. + In the bottom panel of Fig. 4..," In the bottom panel of Fig. \ref{shape}," + we show density maps for all subhaloes belonging to the final haloes of €A2 (Ieft-hand panel) and GASH (right-hand xuiel)., we show density maps for all subhaloes belonging to the final haloes of GA2 (left-hand panel) and GA3n (right-hand panel). + Although these plots are qualitatively similar. there is no detailed: correspondance between subhaloes.," Although these plots are qualitatively similar, there is no detailed correspondance between subhaloes." + On the other hand. the upper panels show that the material which makes up these subhaloes is very similarly distributed. in the two simulations at early epochs.," On the other hand, the upper panels show that the material which makes up these subhaloes is very similarly distributed in the two simulations at early epochs." + Phe biggest dilferences are due to subhaloes which are included in the final halo in one of the simulations but are just outside it in the other., The biggest differences are due to subhaloes which are included in the final halo in one of the simulations but are just outside it in the other. + Fortunately. we do not. care much about the positions of individual subhaloes and are more interested. in statistical results.," Fortunately, we do not care much about the positions of individual subhaloes and are more interested in statistical results." + Aresimulation of an object with higher resolution may not reproduce its structure in detail. but it can still be viewed as the result of evolution from a nearby set of initial conditions (e.g. Llaves 2003).," A re–simulation of an object with higher resolution may not reproduce its structure in detail, but it can still be viewed as the result of evolution from a nearby set of initial conditions (e.g. Hayes 2003)." + A number of authors have arguedIn] that. the. statistical properties of subhaloes in a galaxy-sizecd halo are simply a scaled version of those in a rich cluster halo (Moore et al., A number of authors have argued that the statistical properties of subhaloes in a galaxy-sized halo are simply a scaled version of those in a rich cluster halo (Moore et al. + 1999: οί White 2001: De Lucia et al., 1999; Helmi White 2001; De Lucia et al. + 2004: Diemand et al., 2004; Diemand et al. + 2004)., 2004). + this is surprising. since it is well known that the merging histories of haloes (and in particular their formation times) vary systematically with mass (Lacey Cole 1993: Navarro. Frenk White 1997).," this is surprising, since it is well known that the merging histories of haloes (and in particular their formation times) vary systematically with mass (Lacey Cole 1993; Navarro, Frenk White 1997)." + One might expect. these differences to result in a systematic dependence of the subhalo population On mass., One might expect these differences to result in a systematic dependence of the subhalo population on mass. +" We celine a. dimensionless subhalo mass. m,= Where Alias is the virial mass of the parent halo defined as spherical region which has 200 times critical density of universe at that time."," We define a dimensionless subhalo mass, $m_{n}=m_{sub}/M_{\rm +halo}$ , where $M_{\rm halo}$ is the virial mass of the parent halo defined as spherical region which has 200 times critical density of universe at that time." + In the upper panels of Fig., In the upper panels of Fig. + 5. we plot subhalo abundance against this normalized mass for three ranges of halo mass inour GIE2 simulation. 3.1025TALL.141075.1M]. 1025.1M.3:10H054AAL. and 3e10555fM.1075TALL).," 5, we plot subhalo abundance against this normalized mass for three ranges of halo mass inour GIF2 simulation, $[3\times10^{14}h^{-1}{\rm +M_\odot},\ 10^{15}h^{-1}{\rm M_\odot}]$, $[10^{14}h^{-1}{\rm +M_\odot},\ 3\times 10^{14}h^{-1}{\rm M_\odot}]$ and $[3\times +10^{13}h^{-1}{\rm M_\odot},\ 10^{14}h^{-1}{\rm M_\odot}]$." + These bins contain 7. 33 and 243 haloes. respectively.," These bins contain 7, 33 and 243 haloes, respectively." + Ln this plot. we also include subhalo abundance functions for CX3n. and [ου our 8 cluster simulations., In this plot we also include subhalo abundance functions for GA3n and for our 8 cluster simulations. + Lf halo populations of cdillering mass were just scaled copies of cach other. these various abundance functions would all agree.," If halo populations of differing mass were just scaled copies of each other, these various abundance functions would all agree." + In fact. however. the dillerential and cumulative normalized. mass functions of Fig.," In fact, however, the differential and cumulative normalized mass functions of Fig." + 5 depend. systematically on halo mass., 5 depend systematically on halo mass. + Phe subhalo abundance in high-mass haloes is clearly higher (at. givensealed subhalo mass) than in low-mass haloes., The subhalo abundance in high-mass haloes is clearly higher (at given subhalo mass) than in low-mass haloes. + The dilference between the rich cluster haloes and the galaxy halo Ci3n is à [actor of 2., The difference between the rich cluster haloes and the galaxy halo GA3n is a factor of 2. + The cluster haloes also clearly have more abundant. subhaloes than the lowest mass haloes in our CUP? simulation., The cluster haloes also clearly have more abundant subhaloes than the lowest mass haloes in our GIF2 simulation. + Our simulation cata agree with semi-analytical modelling by Zentner Bullock (2003)., Our simulation data agree with semi-analytical modelling by Zentner Bullock (2003). + These authors argued that. on average. the subhalo mass fraction should. increase with halo mass because high. mass haloes were assembled: more recently.," These authors argued that, on average, the subhalo mass fraction should increase with halo mass because high mass haloes were assembled more recently." + A trend in this direction is also clearly present in in high resolution simulation data of Diemand et al. (, A trend in this direction is also clearly present in in high resolution simulation data of Diemand et al. ( +2004). although these authors emphasise the similarity in subhalo abundance between cluster and galaxy halo rather than the cilference.,"2004), although these authors emphasise the similarity in subhalo abundance between cluster and galaxy halo rather than the difference." + In the bottom panels of Fig., In the bottom panels of Fig. + 5. we show cdillerential ancl cumulative plots of subhalo mass abundance: using a clillerent normalization procedure.," 5, we show differential and cumulative plots of subhalo mass abundance using a different normalization procedure." + We divide the total number of subhaloes in each bin by the total mass of all the parent haloes to obtain the subhalo abundance per unit parent halo mass., We divide the total number of subhaloes in each bin by the total mass of all the parent haloes to obtain the subhalo abundance per unit parent halo mass. + We then plot this abundance as a function of the actual mass (rather than the scaled mass)., We then plot this abundance as a function of the actual mass (rather than the scaled mass). + With this normalization. the subhalo mass functions of dillerent mass haloes agree very well (see also Ixravtsov ct al.," With this normalization, the subhalo mass functions of different mass haloes agree very well (see also Kravtsov et al." + 20042)., 2004a). + For relatively low-mass parent haloes the subhalo abundance drops below that seen in more massive parent haloes for subhalo masses exceeding about 1 per cent of the parent mass., For relatively low-mass parent haloes the subhalo abundance drops below that seen in more massive parent haloes for subhalo masses exceeding about 1 per cent of the parent mass. + lenoring this high mass cut-oll. the subhalo abundance per unit halo mass in Fig.," Ignoring this high mass cut-off, the subhalo abundance per unit halo mass in Fig." + 2 is reasonably well fit by: An immediate consequence of the universality of this relation is a shift with parent halo mass in the abundance of subhalocs as a function of scaled mass m., 2 is reasonably well fit by: An immediate consequence of the universality of this relation is a shift with parent halo mass in the abundance of subhaloes as a function of scaled mass $m_n$ . + For small subhalo masses this shift is, For small subhalo masses this shift is +emissivity associated with oxygen-rich knots trace the position of the reverse shock.,emissivity associated with oxygen-rich knots trace the position of the reverse shock. +" In Fig. 5,,"," In Fig. \ref{fig:xray}," +" we compare the optical [O ΠΠ \ 5007 and Hf emission to a recent X-ray observation of SNR. N132D. On the left plot, we show the X-ray data from Xiao&Chen(2008) (Fig."," we compare the optical [O III] $\lambda$ 5007 and $\beta$ emission to a recent X-ray observation of SNR N132D. On the left plot, we show the X-ray data from \cite{Xiao08} (Fig." + 2 in their paper)., 2 in their paper). + Strong X-ray emission is displayed in purple while weaker emission is in yellow., Strong X-ray emission is displayed in purple while weaker emission is in yellow. +" The characteristic shape of the outer ellipse of emission, reminiscent of a horseshoe, oriented in a NE-SW direction is easily seen."," The characteristic shape of the outer ellipse of emission, reminiscent of a horseshoe, oriented in a NE-SW direction is easily seen." +" On the right hand panel, Fig. 4-"," On the right hand panel, Fig. \ref{fig:above}-" +"-Right, we overlay the X-ray emission over the [ο III] A 5007 and Hf emission."," -Right, we overlay the X-ray emission over the [O III] $\lambda$ 5007 and $\beta$ emission." + Several observations can be made., Several observations can be made. +" First, increased X-ray emissivity along the outer NW edge of the remnant is associated with enhanced optical emission."," First, increased X-ray emissivity along the outer NW edge of the remnant is associated with enhanced optical emission." + This feature is associated to the position of the forward (blast-wave) shock and shocked ISM (Morseetal. 1996)., This feature is associated to the position of the forward (blast-wave) shock and shocked ISM \citep[][]{Morse96}. +. The optical enhancement shows that the ISM shocks are already becoming radiative., The optical enhancement shows that the ISM shocks are already becoming radiative. +" A similar correlation between X-ray and optical emission is also seen in the southern rim, but this region is not completely covered by our observations."," A similar correlation between X-ray and optical emission is also seen in the southern rim, but this region is not completely covered by our observations." +" However, we do see two faint optical knots in the SE portion of the outer blast-wave region, as well as the Runaway Knot (RK), identified previously in Fig. 4.."," However, we do see two faint optical knots in the SE portion of the outer blast-wave region, as well as the Runaway Knot (RK), identified previously in Fig. \ref{fig:above}." +" The abundance of oxygen and lack of Hf emission, as well as its high radial velocity (~ 1000 km 1) does suggest that it is a O-rich clump."," The abundance of oxygen and lack of $\beta$ emission, as well as its high radial velocity $\sim$ 1000 km $^{-1}$ ) does suggest that it is a O-rich clump." + This clump is also associated with a large local enhancement in the X-ray emission., This clump is also associated with a large local enhancement in the X-ray emission. +" However, its position, away from any other oxygen-rich"," However, its position, away from any other oxygen-rich" +"An alternative to leptonic models are (he so-called “hadronie models"" proposed to explain 5-rav enission from blazars.","An alternative to leptonic models are the so-called ""hadronic models"" proposed to explain $\gamma$ -ray emission from blazars." + While leptonic models deal with a relativistic e= plasma in the jet. in hadronic models (he relativistic jet. consists of a relativistic proton (p) and electron (e...) component.," While leptonic models deal with a relativistic $^\pm$ plasma in the jet, in hadronic models the relativistic jet consists of a relativistic proton $p$ ) and electron $e^-$ ) component." + Here we use the hadronic Synchrotron-Proton Dlazar (SPB-) model of Mückeetal.(2002). (o model the spectral energy. distribution (SED) of Wo Comae in Mav 1993., Here we use the hadronic Synchrotron-Proton Blazar (SPB-) model of \cite{muecke02} to model the spectral energy distribution (SED) of W Comae in May 1998. + Like in the leptonic model (he emission region. or “blob”. in an AGN jet moves relativistically along the jet axis which is closely aligned with our line-ol-sight.," Like in the leptonic model the emission region, or “blob”, in an AGN jet moves relativistically along the jet axis which is closely aligned with our line-of-sight." +" Relativistic (accelerated) protons. whose particle density πρ follows a power law spectrum x5,an"" in the range 2< μπας. ave Injected instantaneously into a highly magnetized environment (2=cons! within (he emission region). and sulfer enerev losses due to protonphoton interactions (meson production aud Bethe-Ieitler pair production). svnchrotron radiation and adiabatic expansion."," Relativistic (accelerated) protons, whose particle density $n_p$ follows a power law spectrum $\propto \gamma_p^{-\alpha_p}$ in the range $2\leq\gamma_p\leq\gamma_{\rm{p,max}}$ , are injected instantaneously into a highly magnetized environment $B=const$ within the emission region), and suffer energy losses due to proton–photon interactions (meson production and Bethe-Heitler pair production), synchrotron radiation and adiabatic expansion." + The mesons produced in photonmeson interactions always decay in astroplivsical environments. however. (hey may sulfer svnchrotron losses before the decay. which is taken into account in this moclel.," The mesons produced in photonmeson interactions always decay in astrophysical environments, however, they may suffer synchrotron losses before the decay, which is taken into account in this model." + The relativistic primary ο radiate svnchirotron photons that manifest themselves in the blazar SED as the svuchrotron hump. and serve as the target radiation field for proton-photon interactions and the pair-svnchrotron cascade which subsequently develops.," The relativistic primary $e^{-}$ radiate synchrotron photons that manifest themselves in the blazar SED as the synchrotron hump, and serve as the target radiation field for proton-photon interactions and the pair-synchrotron cascade which subsequently develops." + The SPD-mocdel is designed for objects with a negligible external target photon component. aud hence suitable or BL Lac Objects.," The SPB-model is designed for objects with a negligible external target photon component, and hence suitable for BL Lac Objects." + The cascade redistributes (he photon power to lower energies where the photons eventually escape from the emission region., The cascade redistributes the photon power to lower energies where the photons eventually escape from the emission region. +" The cascades can be initiated by photons rom z""-decav (7x"" cascade?). electrons from the w=—p=€7 decay (737 cascade?). p- photons (p-synchrotron cascade""). charged je. π- and A-svnchrotron photons (""pu7-8ynehrotron cascade"") and ο from the proton-photon Dethe-IHHeitller pair production (""Dethe-LHeitler cascade?)."," The cascades can be initiated by photons from $\pi^0$ -decay $\pi^0$ cascade”), electrons from the $\pi^\pm\to \mu^\pm\to e^\pm$ decay $\pi^\pm$ cascade”), $p$ -synchrotron photons $p$ -synchrotron cascade”), charged $\mu$ -, $\pi$ - and $K$ -synchrotron photons $\mu^\pm$ -synchrotron cascade”) and $e^\pm$ from the proton-photon Bethe-Heitler pair production (“Bethe-Heitler cascade”)." +" Mücke&Protheroe(2001). and Mückeetal.(2002) have shown that the 7x"" cascades? and πι cascades” generate rather featureless photon spectra. in contrast (ο 7p-synchrotron cascades? and 7j7-synchrotron cascades? that produce a double-humped SED as typically observed. for 5-rav blazars."," \cite{mp01} and \cite{muecke02} have shown that the $\pi^0$ cascades” and $\pi^\pm$ cascades” generate rather featureless photon spectra, in contrast to $p$ -synchrotron cascades” and $\mu^\pm$ -synchrotron cascades” that produce a double-humped SED as typically observed for $\gamma$ -ray blazars." + The contribution from the Dethe-IHeitler cascades is mostly negligible., The contribution from the Bethe-Heitler cascades is mostly negligible. + In general direct proton and muon svnchrotron radiation is mainly responsible for (he hieh energv hump in blazars whereas the low energy hump is dominanted bv svnchrotron radiation [rom the primary €... wilh a contribution of svnchrotron radiation [rom secondary electrons (produced bvthe p- and. jc7-synchrotron cascade).," In general direct proton and muon synchrotron radiation is mainly responsible for the high energy hump in blazars whereas the low energy hump is dominanted by synchrotron radiation from the primary $e^-$, with a contribution of synchrotron radiation from secondary electrons (produced bythe $p$ - and $\mu^\pm$ -synchrotron cascade)." +" A detailed description of the model itself, ancl its implementation as a (time-independent) Monte-Carlo/nunmerical code. has been given in Mücke&Protheroe (2001).."," A detailed description of the model itself, and its implementation as a (time-independent) Monte-Carlo/numerical code, has been given in \cite{mp01}. ." +slopes. based ou the OGLE-IHI Cepheids. became comparaje to the slopes give1 iu wlen only the long period (log[P?]71) Cepheids were used.,"slopes, based on the OGLE-III Cepheids, became comparable to the slopes given in \citet{mad09} when only the long period $\log[P]>1$ ) Cepheids were used." + For tle GalacticAC oiuxd P-L relatious. the P-L sopes derived from tle infrared surface rightuess (IRSB) techuique (CALI and GAL? in Table 2) ) were steeper than he slopes derived rou Cepleids that have paallax. Measurements (CG:AL3 i1 Table 3)) .," For the Galactic band P-L relations, the P-L slopes derived from the infrared surface brightness (IRSB) technique (GAL1 and GAL2 in Table \ref{slopecompare}) ) were steeper than the slopes derived from Cepheids that have parallax measurements (GAL3 in Table \ref{slopecompare}) )." + The laree uncertainties of he GAL3 P-L slopes (dueothecombinationofasmallnumberdatapoiusinsampleaudlessp'ecisephotometry.seeMarengoeal.2010) Calse all tliree sets of Galactic P-L slopes o be cousistent with each oher.," The large uncertainties of the GAL3 P-L slopes \citep[due to the combination of a small number of data points in the sample and less precise photometry, see][]{mar10} cause all three sets of Galactic P-L slopes to be consistent with each other." + It is wortwhile to uote tiat the theoretical P-L slopes derived in Marengoeal.(2010) agree well witl he GAL3 P-L slopes., It is worthwhile to note that the theoretical P-L slopes derived in \citet{mar10} agree well with the GAL3 P-L slopes. + If the uncertaiules of the P-L slopes derived frou the parallax clistauces :ure the same as hose from the IRSB echuiques. theu he slopes betweei them. will disagree at tle 3.50 level.," If the uncertainties of the P-L slopes derived from the parallax distances are the same as those from the IRSB techniques, then the slopes between them will disagree at the $\sim 3.5\sigma$ level." + lareugoetal.(2010 suggestedMOD that he projection [acor-period relation used i IRSB techuique may need further reineinent., \citet{mar10} suggested that the projection factor-period relation used in IRSB technique may need further refinement. + Finally. we liwe derived the? baud (at ~Bye1) P-L relation for the LMC Cepheils. which is incepeneeut of the SAGE data. u Neeowetal.(2010).," Finally, we have derived the band (at $\sim3\mu{\mathrm m}$ ) P-L relation for the LMC Cepheids, which is independent of the SAGE data, in \citet{nge10}." +. The slope of te N3 baud P-L relation is —3.2540.05., The slope of the $N3$ band P-L relation is $-3.25\pm0.05$. + This agrees well wih the 3.65 batcL P-L slopes in Table 2 for the slopes in the “Shallow $ope group., This agrees well with the $3.6\mu{\mathrm m}$ band P-L slopes in Table \ref{slopecompare} for the slopes in the “Shallow Slope” group. + Iu this Paper we have derived SMC P-L relations in the bauds. by matching the archival SACE-SMC data to the latest SMC Cepheid catalog from OCLE-ILL.," In this Paper we have derived SMC P-L relations in the bands, by matching the archival SAGE-SMC data to the latest SMC Cepheid catalog from OGLE-III." + We lave fouud. for the first time. that the SNC P-L 'elations show a change of slope at log(/?)=0.1 in the 3.6740 and. L5 baud. similar to tel optical counterparts.," We have found, for the first time, that the SMC P-L relations show a change of slope at $\log(P)=0.4$ in the $3.6\mu{\mathrm m}$ and $4.5\mu{\mathrm m}$ band, similar to their optical counterparts." + Due to the incompleteness at the shor period eud. such athange of slope cau101 ye coulirmed to exist in the 5.8421 and 8.0421 bands.," Due to the incompleteness at the short period end, such a change of slope cannot be confirmed to exist in the $5.8\mu{\mathrm m}$ and $8.0\mu{\mathrm m}$ bands." + Fiture observations wibhJWST may elp t« determine the change of slope in these two bauds., Future observations with may help to determine the change of slope in these two bands. + The slopes of tl esIC.IRAC baud P-L relatious were found to be around —23.2. which is sistent with te LAC ‘slopes found in Ngeowetal.(2009b) aud the Calactie P-L relations derived [rom Cepeic SVIz‘ith parallax measurements (Marengoetal.2010).," The slopes of the SMC band P-L relations were found to be around $-3.2$, which is consistent with the LMC slopes found in \citet{nge09} and the Galactic P-L relations derived from Cepheids with parallax measurements \citep{mar10}." +. This judicates that the slopes of the P-L relatio isin theJRAC bauds are inseusitive to metallicity., This indicates that the slopes of the P-L relations in the bands are insensitive to metallicity. + However. the SMC slopes disagree wih he steeper slopes for the LMC and Galactic counterparts from and those from the IRSB technique. respectively.," However, the SMC slopes disagree with the steeper slopes for the LMC and Galactic counterparts from \citet{mad09} and those from the IRSB technique, respectively." + Future observations such as the parallax ineasurements [rouCare for Galactic Cepheids and/orJIWST observations of Magellanic Cloud Cepheids may hel» to resolve this discrepancy., Future observations such as the parallax measurements from for Galactic Cepheids and/or observations of Magellanic Cloud Cepheids may help to resolve this discrepancy. + CCN thanks the fuucdiug from National Scieuce Council (of Taiwan) under the contract. NSC 98-2112-M-008-013-MY3., CCN thanks the funding from National Science Council (of Taiwan) under the contract NSC 98-2112-M-008-013-MY3. + We would like to thauk the referee for helpful comments to improve the manuscript., We would like to thank the referee for helpful comments to improve the manuscript. + We would also like to thank Hikliug Neilson aud Nauey Evans for useful discussions., We would also like to thank Hilding Neilson and Nancy Evans for useful discussions. + This work is based [in part] on observatious made with the Spitzer Space Telescope. which is," This work is based [in part] on observations made with the Spitzer Space Telescope, which is" +"The interstellar medium in the dmunediate solar neighborhood is part of the ""local interstellar cloud” (LIC).",The interstellar medium in the immediate solar neighborhood is part of the “local interstellar cloud” (LIC). + The flow of the partially ionized LIC past the Suu constitutes a pressure that balances aud terminates the expansion of the coronal solar wind., The flow of the partially ionized LIC past the Sun constitutes a pressure that balances and terminates the expansion of the coronal solar wind. + These two wiuds create a morphology that includes the termination shock rausition of the supersouic solar wind to a hot wcliosheatl or heliotzil flow., These two winds create a morphology that includes the termination shock transition of the supersonic solar wind to a hot heliosheath or heliotail flow. + Au interstellar bow shock is likely to be recessary as well o decelerate the LIC dow., An interstellar bow shock is likely to be necessary as well to decelerate the LIC flow. + The ionized Hows of the LIC aud the solar wine are separated by ie heliopause., The ionized flows of the LIC and the solar wind are separated by the heliopause. + The LIC also supplies the svsteia with oeiterstellar neutrals. predominantly with neutral hydrogen (II).," The LIC also supplies the system with interstellar neutrals, predominantly with neutral hydrogen (H)." + Neutral II interacts weakly with the plasma. mainly woueh charec exchanges with plasma protons.," Neutral H interacts weakly with the plasma, mainly through charge exchanges with plasma protons." + The distance even to the termination shock is large rough that there are only a few iu-tu iueasurenments o date in the outer heliosplere., The distance even to the termination shock is large enough that there are only a few in-situ measurements to date in the outer heliosphere. +" Notable sources of information are the two Vovager spacecraft at a distance of LOL AU and 51 AU from the Sun (2007 September). respectively, with Vovaecr having assed into the heliosheath region ou 2001 December 16 (e.g.Stoneetal.Ww 105).. and Voyager 2 on 2007 August 30."," Notable sources of information are the two Voyager spacecraft at a distance of 104 AU and 84 AU from the Sun (2007 September), respectively, with Voyager 1 having passed into the heliosheath region on 2004 December 16 \citep[e.g., ][]{Stone05}, and Voyager 2 on 2007 August 30." + For examples of in-depth analyses of observations relating to the outer heliosphere. we refer to other contributions iu this specia issue (Dzowskietal.2008:PrvorRichardson(al.2008:Slavin&Frisch 2008).," For examples of in-depth analyses of observations relating to the outer heliosphere, we refer to other contributions in this special issue \citep{Bzowski08,Pryor08,Richardson08,Slavin08}." +. Data from the outer heliosphere are sparse. iux uunuercal modoeliug of the elobal heliosphere/LIC. svsteii plavs au iuiportant role for the analvsis aud iuterpretation of observations.," Data from the outer heliosphere are sparse, and numerical modeling of the global heliosphere/LIC system plays an important role for the analysis and interpretation of observations." + It is needed to relate the uudisturbe LIC dow and its physical parameters to the processed and changed flow that we observe im the heliosphere inside the termunation shock., It is needed to relate the undisturbed LIC flow and its physical parameters to the processed and changed flow that we observe in the heliosphere inside the termination shock. + Iu fundamental wavs all the LIC constraints formulated in the accompanying papers (Dzowskietal.2008:PrvorctRichardson2008:Slawin&Frisch2008) involve elobal heliosphere modcling.," In fundamental ways all the LIC constraints formulated in the accompanying papers \citep{Bzowski08,Pryor08,Richardson08,Slavin08} involve global heliosphere modeling." + Also the evaluation of future data sets from the Iuterstellay Boundary Explorer (IBEX) mission. which focuses on secondary ueutrals created in the heliosphere and on the LIC oxygen and helimu flow through tli heliosphere (AleComasetal.200L).. depends crucially on this kind of modeling.," Also the evaluation of future data sets from the Interstellar Boundary Explorer (IBEX) mission, which focuses on secondary neutrals created in the heliosphere and on the LIC oxygen and helium flow through the heliosphere \citep{IBEX}, depends crucially on this kind of modeling." + All such elobal models imake asstmptious aud sinplificatious. most often with the goal of isolating the influcuce of a specific plivsical effect. (c.g... the tilt of the LIC inagnetie field with respect to the LIC flow vector). or in order to keep computation times reasonable.," All such global models make assumptions and simplifications, most often with the goal of isolating the influence of a specific physical effect (e.g., the tilt of the LIC magnetic field with respect to the LIC flow vector), or in order to keep computation times reasonable." + The identification of hehospheric asvuuuectrics with respect to the helimu LIC flow vector (Móbiusetal.2001:Lalle-ientetal.2005.andreferencestherein) has increase interest iu the development of realistic. three-«dinieusiona (3D) MIID models. as differeut oricutious and strenetls of the interstellar maenetic field can help to explain these asvuuuetrics.," The identification of heliospheric asymmetries with respect to the helium LIC flow vector \citep[][and references +therein]{Moebius04,Lallement05} has increased interest in the development of realistic, three-dimensional (3D) MHD models, as different orientations and strengths of the interstellar magnetic field can help to explain these asymmetries." + However. the fact remains that Letra interstellar II cutering the heliosphere has a more decisive oeiduence on the helospheric shape. extent. and particle content.," However, the fact remains that neutral interstellar H entering the heliosphere has a more decisive influence on the heliospheric shape, extent, and particle content." + For this reason. we focus here on nuuerica models that treat the plasmaneutral iuteractiou in a sclconsistent wav. but ucelect the influence of iuterplauetary or interstellar magnetic fields.," For this reason, we focus here on numerical models that treat the plasma/neutral interaction in a self-consistent way, but neglect the influence of interplanetary or interstellar magnetic fields." + The models are. iu principle. 3D) plasma/ueutral codes for which plasma and neutrals are coupled by charee exchauge.," The models are, in principle, 3D plasma/neutral codes for which plasma and neutrals are coupled by charge exchange." + Wherever it is possible. the assumption of azimuthal svuuuctry," Wherever it is possible, the assumption of azimuthal symmetry" + , +ft that relates ¢ to the gr andr2 colors: Equation 3 is valid over the ranges shown in Fieure 10 aud has typical uncertainties of,fit that relates $\zeta$ to the $g-r$ and $r-z$ colors: Equation 3 is valid over the ranges shown in Figure \ref{zetacol} and has typical uncertainties of. + Prompted by the discovery of new low-nmiass stars in the solar vicinity from other SDSS studies (e.@.Schiaidtetal. 2010a).. we inspected our catalog for AL dwarts that are possibly within 25 pe of the Sun.," Prompted by the discovery of new low-mass stars in the solar vicinity from other SDSS studies \citep[e.g.][]{schmidt_blue}, we inspected our catalog for M dwarfs that are possibly within 25 pc of the Sun." + We found 21 nearby M chwart caucidates (8 of which were previously nuidentified) based ou the distances derived using the Myr0i photometric parallax. relation of Bochauskietal. (2010).," We found 21 nearby M dwarf candidates (8 of which were previously unidentified) based on the distances derived using the $M_r$, $r-z$ photometric parallax relation of \citet{boo10}." +. Three of the stars were in the WWtüs SDSS DR5 sample but were not identified there as potential solar neighbors., Three of the stars were in the \nocite{west08}W W08 SDSS DR5 sample but were not identified there as potential solar neighbors. + Table 5 gives the positions. spectral types and distance estimates for all 21 of the nearby candidates.," Table \ref{table:dist} gives the positions, spectral types and distance estimates for all 21 of the nearby candidates." + The closest candidate is SDSSI110|15816. which has au estimated distance of 118 pc.," The closest candidate is SDSS1410+1846, which has an estimated distance of 14.8 pc." + Since roughly oue out of every 1000 AL chwarts in SDSS has a spectra. there is a strong possibility that many more nearby stars await discovery and will be cataloged iu future stucies.," Since roughly one out of every 1000 M dwarfs in SDSS has a spectrum, there is a strong possibility that many more nearby stars await discovery and will be cataloged in future studies." + We have presented the SDSS DR? M dwarf spectroscopic catalog. which consists of more than TO.OQOQ. visually confirmed AL dwarfs aud represcuts the largest spectroscopic sample of M dwarfs ever asseaubled.," We have presented the SDSS DR7 M dwarf spectroscopic catalog, which consists of more than 70,000 visually confirmed M dwarfs and represents the largest spectroscopic sample of M dwarfs ever assembled." + Om value-added catalog includes proper motions. RVs. photometric matches to 2MLASS. spectral classification. distances. calculated space velocities. activitv-1uduced enüssiou liue measurements and molecular baudblead strengths.," Our value-added catalog includes proper motions, RVs, photometric matches to 2MASS, spectral classification, distances, calculated space velocities, activity-induced emission line measurements and molecular bandhead strengths." + The DR? catalog is available for download at the Vizier site or bv contacting the correspondiug author., The DR7 catalog is available for download at the Vizier site or by contacting the corresponding author. + Our analysis of the visual spectral classification as compared to the automatic Wamuuer results reveals a slight ~0.1 subtype svstematic offset in the automatic ILhunnmer spectral types for late-type ΑΙ cawarts but confirms that the automatic types are good to within the stated precision of l spectral subtype., Our analysis of the visual spectral classification as compared to the automatic Hammer results reveals a slight $\sim0.4$ subtype systematic offset in the automatic Hammer spectral types for late-type M dwarfs but confirms that the automatic types are good to within the stated precision of 1 spectral subtype. + We present updated median colors for all M dwarf spectral types., We present updated median colors for all M dwarf spectral types. + We lave also analyzed some of the bulk properties of the DR sample., We have also analyzed some of the bulk properties of the DR7 sample. + The main results of our analysis are as follows: Many additional studies will make use of our AL dwarf catalog., The main results of our analysis are as follows: Many additional studies will make use of our M dwarf catalog. + Those already underway inchide a statistical, Those already underway include a statistical +pe. and curiously this is not at all far from the median distance of each source to its nearest neighbor: ~1.8pe for the ssources In the /=59° ffield for the sources” average distance (Russeiletal..20100.,"pc, and curiously this is not at all far from the median distance of each source to its nearest neighbor: $\sim$ 1.8pc for the sources in the $l$ field for the sources' average distance \citep{russeil10}." + It is also interesting that the typical fragmentation length scale decreases with increasing filament density. in broad qualitative agreement with a higher spatial density of sources in the brightest filaments.," It is also interesting that the typical fragmentation length scale decreases with increasing filament density, in broad qualitative agreement with a higher spatial density of sources in the brightest filaments." + This clearly deserves further investigation to be confirmed as a viable hypothesis., This clearly deserves further investigation to be confirmed as a viable hypothesis. + The first science highlights presented in this paper. as well as in the accompanying papers in this volume and elsewhere. show that owing to its optimal use of unique characteristics of wavelength coverage. spatial resolution and mapping speed. the ssurvey has the potential to lead to a quantum leap in our understanding of large-scale Galactie star formation from cloud to cluster-forming clump formation and of the evolution of protoclusters and massive protostars.," The first science highlights presented in this paper, as well as in the accompanying papers in this volume and elsewhere, show that owing to its optimal use of unique characteristics of wavelength coverage, spatial resolution and mapping speed, the survey has the potential to lead to a quantum leap in our understanding of large-scale Galactic star formation from cloud to cluster-forming clump formation and of the evolution of protoclusters and massive protostars." +FWHM of 36 mas at 10 um. This interpretation suggests that. apart from the central emission. about of the total flux is (smoothly) radiated from (projected) distances greater than 150 AU from the star. (,"FWHM of 36 mas at 10 $\mu$ m. This interpretation suggests that, apart from the central emission, about of the total flux is (smoothly) radiated from (projected) distances greater than 150 AU from the star. (" +2) Another common interpretation aims on estimating an average angular scale-length of the entire flux distribution by fitting a Gaussian profile to the observed visibility measurements.,2) Another common interpretation aims on estimating an average angular scale-length of the entire flux distribution by fitting a Gaussian profile to the observed visibility measurements. + The measured flat visibility spectrum results in a wavelength-dependent FWHM: 24. 30. and 39 mas at 8. 10. and 13 jam. respectively.," The measured flat visibility spectrum results in a wavelength-dependent FWHM: 24, 30, and 39 mas at 8, 10, and 13 $\mu$ m, respectively." + An increasing diameter towards longer wavelengths is expected for optically thin. thermal. sources showing cooler temperatures. which dominate the flux at longer wavelengths. further out.," An increasing diameter towards longer wavelengths is expected for optically thin, thermal sources showing cooler temperatures, which dominate the flux at longer wavelengths, further out." + Summarizing. the MIR-brightness distribution of IRS 7 is clearly dominated by dust emission extended on 30 mas scales and greater.," Summarizing, the MIR-brightness distribution of IRS 7 is clearly dominated by dust emission extended on 30 mas scales and greater." + Accordingly. the fitted black-body temperature of only 200 K (Fig. 2))," Accordingly, the fitted black-body temperature of only 200 K (Fig. \ref{fig:2}) )" + appears to fit the complete N-band emission., appears to fit the complete $N$ -band emission. + This (for close stellar environments) relatively low dust temperature is another confirmation. of observing dust relatively far away from the central heat source., This (for close stellar environments) relatively low dust temperature is another confirmation of observing dust relatively far away from the central heat source. + A possible explanation of the origin of this extended dust emission (and of the low visibilities) derives from radio wavelengths., A possible explanation of the origin of this extended dust emission (and of the low visibilities) derives from radio wavelengths. +" More than 15 years ago ?. spatially resolved in 2 em VLA maps of 0.4 "" angular resolution a bow-shock-like free-free radio emission region around IRS 7.", More than 15 years ago \citet{1991ApJ...371L..59Y} spatially resolved in 2 cm VLA maps of 0.4 $\arcsec$ angular resolution a bow-shock-like free-free radio emission region around IRS 7. + The curved emission morphology shields IRS 7. which Hes ~ 150 mas North of the apex of the shock front according to the model of ?.. from the heating and tonizing galactic winds emanating from the very center of the Galaxy.," The curved emission morphology shields IRS 7, which lies $\sim$ 150 mas North of the apex of the shock front according to the model of \citet{1992ApJ...385L..41Y}, from the heating and ionizing galactic winds emanating from the very center of the Galaxy." + It is conceivable that (a part of) the extended MIR emission detected by MIDI is generated in this externally heated environment., It is conceivable that (a part of) the extended MIR emission detected by MIDI is generated in this externally heated environment. + The radiative dominance of this emission over the warmer dust closer to the star might point to the outer dust either being optically thick. or overwhelming the inner. slightly hotter dust closer to IRS 7 due to an energetic domination of the external radiation field over heating by IRS 7 itself.," The radiative dominance of this emission over the warmer dust closer to the star might point to the outer dust either being optically thick, or overwhelming the inner, slightly hotter dust closer to IRS 7 due to an energetic domination of the external radiation field over heating by IRS 7 itself." + In line with this interpretation is that the position angle of the projected baseline during the time of observation is with PAx30° (East-of-North) relatively close to the roughly North-South elongation of the radio bow-shock.," In line with this interpretation is that the position angle of the projected baseline during the time of observation is with $PA\,\approx\,30^\circ$ (East-of-North) relatively close to the roughly North-South elongation of the radio bow-shock." + This would explain the low amount of correlated MIR-flux., This would explain the low amount of correlated MIR-flux. + A better «-coverage is needed to prove this scenario., A better $uv$ -coverage is needed to prove this scenario. + Estimating the putative apex and radiative properties of a MIR shock front around IRS 7 pointing to the GC will allow us to estimate physical properties complementing the radio data. and to investigate the interstellar conditions only a few are-seconds away from the MBH.," Estimating the putative apex and radiative properties of a MIR shock front around IRS 7 pointing to the GC will allow us to estimate physical properties complementing the radio data, and to investigate the interstellar conditions only a few arc-seconds away from the MBH." + The calibrated total-flux spectrum from. the MIDI photometry of IRS 7 shows a broad interstellar jam-silicate absorption feature.," The calibrated total-flux spectrum from the MIDI photometry of IRS 7 shows a broad interstellar $\,\mu$ m-silicate absorption feature." + The optical depth tosjn *7 im this feature is twice as high as expected for the usually adopted interstellar MIR. extinction. towards the central parsee. as derived from earlier observations of the GC at much lower angular resolution.," The optical depth $\tau_{\rm 9.8~\mu m}\,\approx\,7$ in this feature is twice as high as expected for the usually adopted interstellar MIR extinction towards the central parsec, as derived from earlier observations of the GC at much lower angular resolution." + This τοςgm resembles our recent findings towards the nearby IRS 3 (?)., This $\tau_{\rm 9.8~\mu m}$ resembles our recent findings towards the nearby IRS 3 \citep{2007arXiv0711.0249P}. +" This strong Tos 4,-eXCess cannot solely be attributed to circumstellar dust enshrouding only these two stars for various reasons.", This strong $\tau_{\rm 9.8~\mu m}$ -excess cannot solely be attributed to circumstellar dust enshrouding only these two stars for various reasons. + The spectral shape of the 9.8 jme-silicate. feature fits the typical shape ofinterstellar silicate dust absorption., The spectral shape of the 9.8 $\mu$ m-silicate feature fits the typical shape of silicate dust absorption. + The NIR-continuum extinction towards these two stars (?) does not show a similar eXCess. and circumstellar silicate dust normally shows the silicate feature in emission.," The NIR-continuum extinction towards these two stars \citep{2007A&A...469..125S} does not show a similar excess, and circumstellar silicate dust normally shows the silicate feature in emission." + This paradox of normal extinction at NIR and relatively enhanced Tuyµη suggests a relative silicate overabundance over carbonaceous dust grains in the diffuse interstellar environment of the entire central parsec. or at least around IRS 3 and IRS 7.," This paradox of normal extinction at NIR and relatively enhanced $\tau_{\rm 9.8~\mu m}$ suggests a relative silicate overabundance over carbonaceous dust grains in the diffuse interstellar environment of the entire central parsec, or at least around IRS 3 and IRS 7." + Further support comes from the MIDI spectro-visibilities of IRS 7., Further support comes from the MIDI spectro-visibilities of IRS 7. + The essentially flat curve means that the dust resolved by the interferometer. and located to the star. doesnof account for a major fraction of the excess of interstellar silicate absorption. since otherwise the change in optical depth along the silicate feature would be visible às à spectral feature in the visibility spectrum.," The essentially flat curve means that the dust resolved by the interferometer, and located to the star, does account for a major fraction of the excess of interstellar silicate absorption, since otherwise the change in optical depth along the silicate feature would be visible as a spectral feature in the visibility spectrum." +" Here again. ""close to the star’ means closer than about 30 mas."," Here again, 'close to the star' means closer than about 30 mas." + The flat visibility over the N-band thus suggests that the dust responsible for the τονjn-excess is located further away from the star. diffusely distributed within the central parsec. and suggests an unusual silicate overabundance in the GC-dust chemistry.," The flat visibility over the $N$ -band thus suggests that the dust responsible for the $\tau_{\rm 9.8~\mu m}$ -excess is located further away from the star, diffusely distributed within the central parsec, and suggests an unusual silicate overabundance in the GC-dust chemistry." + That the. usually linear. relation between the optical depth in the silicate feature and the extinction at other wavelengths. known from various studies and locations. throughout the galaxy. can break. down in regions with particular properties has recently been confirmed for dense interstellar clouds (?)..," That the, usually linear, relation between the optical depth in the silicate feature and the extinction at other wavelengths, known from various studies and locations throughout the galaxy, can break down in regions with particular properties has recently been confirmed for dense interstellar clouds \citep{2007ApJ...666L..73C}." + In the Galactic. center. recent observations report the possibility of an unusual stellar initial mass function (e.g.??).. which could affect the GC dust chemistry.," In the Galactic center, recent observations report the possibility of an unusual stellar initial mass function \citep[e.g.][]{2007MNRAS.374L..29K,2007ApJ...669.1024M}, which could affect the GC dust chemistry." + Also ?— derived. from. chemical abundance measurements of 10 GC stellar atmospheres within the central 30 pe enhanced [O/Fe| and [Ca/Fe] ratios. and concluded an overabundance of a-elements in the GC. including Si.," Also \citet{2007ApJ...669.1011C} derived from chemical abundance measurements of 10 GC stellar atmospheres within the central 30 pc enhanced [O/Fe] and [Ca/Fe] ratios, and concluded an overabundance of $\alpha$ -elements in the GC, including Si." +" This could explain an interstellar Tos ,4,-excess with respect to the carbon-dominated NIR-extinetion. as suggested by our MIDI data."," This could explain an interstellar $\tau_{\rm 9.8~\mu m}$ -excess with respect to the carbon-dominated NIR-extinction, as suggested by our MIDI data." + We present the first successful NIR fringe measurements with an optical long baseline interferometer of a star in the Galactic center at a nominal angular resolution of 9 mas., We present the first successful NIR fringe measurements with an optical long baseline interferometer of a star in the Galactic center at a nominal angular resolution of 9 mas. + The analysis of the 50 m-baseline K-band visibilities of the red supergiant leads to a uniform-dise diameter of 2.6 mas which suggests that all the VLTI-UT baselines up to a length of 130 m resolve the starlight., The analysis of the 50 m-baseline $K$ -band visibilities of the red supergiant leads to a uniform-disc diameter of 2.6 mas which suggests that all the VLTI-UT baselines up to a length of 130 m resolve the starlight. + This size is larger than the expected diameter of the photosphere and suggests the detection of hot circumstellar dust and molecular shells., This size is larger than the expected diameter of the photosphere and suggests the detection of hot circumstellar dust and molecular shells. + But more NIR-data at high SNR are needed to confirm this result. and to estimate an accurate diameter of IRS 7.," But more NIR-data at high SNR are needed to confirm this result, and to estimate an accurate diameter of IRS 7." + The strong interferometric signal observed shows the feasibility of future interferometric experiment involving IRS 7., The strong interferometric signal observed shows the feasibility of future interferometric phase-referencing experiment involving IRS 7. + First MIR interferometric data taken on IRS 7 with the VLTI-MIDI instrument proved that most of the flux at 10 jum , First MIR interferometric data taken on IRS 7 with the VLTI-MIDI instrument proved that most of the flux at 10 $\mu$ +nol observed.,not observed. + Furthermore. it contradiets previous findings which suggest that ILDEG are. in anv case. older than LDEG (Ixuntschueretal.2002)).," Furthermore, it contradicts previous findings which suggest that HDEG are, in any case, older than LDEG \citealt{Kun02}) )." + Thus. the most plausible explanation of our results is that variations in the relative abundances of C and N with respect to Mg and Fe are responsible for the observed olfsets between galaxies in different environments.," Thus, the most plausible explanation of our results is that variations in the relative abundances of C and N with respect to Mg and Fe are responsible for the observed offsets between galaxies in different environments." + Fig., Fig. +" 3. compares the observed CN»à and C4668"" with the predictions of stellar population models.", \ref{index-index} compares the observed $_{2}$ and $'$ with the predictions of stellar population models. + This ligure clearly shows that (he previously found overabundances of C and N only stand for LDEG. while IIDEG tend to exhibit relative abundances closer to the solar partition.," This figure clearly shows that the previously found overabundances of C and N only stand for LDEG, while HDEG tend to exhibit relative abundances closer to the solar partition." + The C4668' index is extremely sensitive to carbon changes. so little variations in carbon abundance can change this index dramatically (Iripicco&Dell 1995)).," The $'$ index is extremely sensitive to carbon changes, so little variations in carbon abundance can change this index dramatically \citealt{Tri95}) )." + However. the variations of CN are mostly controlled by N. because extra C is readily incorporated into CO but extra N makes more CN molecules.," However, the variations of CN are mostly controlled by N, because extra C is readily incorporated into CO but extra N makes more CN molecules." + Therefore. a change in both C and N abundances is required {ο explain our results.," Therefore, a change in both C and N abundances is required to explain our results." + Desides. il extra carbon is easily incorporated into CO. one should detect an enhancement in (he strength of the CO bands in LDEG compared to IIDEG.," Besides, if extra carbon is easily incorporated into CO, one should detect an enhancement in the strength of the CO bands in LDEG compared to HDEG." + This effect has indeed been found by Mobasher&James(1996) comparing the CO band at 2.3j m in ealaxies from the field and from the Pisces ancl Abell 2634 clusters.," This effect has indeed been found by \citet{Mob96} + comparing the CO band at $\mu$ m in galaxies from the field and from the Pisces and Abell 2634 clusters." + They interpreted (his difference as an evidence of intermecdiate-age stellar population in LDEG (through a major contribution of AGB stars)., They interpreted this difference as an evidence of intermediate-age stellar population in LDEG (through a major contribution of AGB stars). + Although we do not discard a larger contribution of vounger populations in LDEG compared to ILDEG. our results in the blue spectral range imply that their observations can be solely explained by a relative abundance difference.," Although we do not discard a larger contribution of younger populations in LDEG compared to HDEG, our results in the blue spectral range imply that their observations can be solely explained by a relative abundance difference." + In any. case. the observed offsets can not be due (to an age effect alone (see above).," In any case, the observed offsets can not be due to an age effect alone (see above)." + The most plausible scenario to explain the differences in relative abundances is (he one in whieh LDEG and ΠΟΙΟ have experienced dilferent star formation histories., The most plausible scenario to explain the differences in relative abundances is the one in which LDEG and HDEG have experienced different star formation histories. + In particular. since the ISM is progressively enriched in C and N by stars of low and intermediate mass stars. DEG should have been fully assembled belore the massive release of these elements.," In particular, since the ISM is progressively enriched in C and N by stars of low and intermediate mass stars, HDEG should have been fully assembled before the massive release of these elements." + The hierarchical clustering paradigm currently. predicts (hat galaxies in clusters formed. αἱ different epochs than. LDEGs., The hierarchical clustering paradigm currently predicts that galaxies in clusters formed at different epochs than LDEGs. + If the time elapsed between (he assembling of the former and the later is enough (o permit the C and N enrichment of the ISA of the pre-mereine building blocks in LDEG. the stars lormed in these galaxies during the merging events will be C and N enhanced.," If the time elapsed between the assembling of the former and the later is enough to permit the C and N enrichment of the ISM of the pre-merging building blocks in LDEG, the stars formed in these galaxies during the merging events will be C and N enhanced." + The constancy of the iron-peak elements could be understood il. in both environments. the mergers (ake place before (vpe-I SN can significantly pollute the ISM of the pre-merging blocks.," The constancy of the iron-peak elements could be understood if, in both environments, the mergers take place before type-I SN can significantly pollute the ISM of the pre-merging blocks." + Additionally. IDEG could have experienced a truncated star ormation and chemical enrichment history. compared to a more continuous time-extended ustory for their counterparts in low density environments.," Additionally, HDEG could have experienced a truncated star formation and chemical enrichment history compared to a more continuous time-extended history for their counterparts in low density environments." + However. under this hypothesis. there shoukl be an increase of magnesium (produced by (vpe-II SN) in LDEG which is 100 detected.," However, under this hypothesis, there should be an increase of magnesium (produced by type-II SN) in LDEG which is not detected." + One way (o understand the constancy of the Meh’ index could be to invoke, One way to understand the constancy of the $'$ index could be to invoke +alihough from (the available data we are not able to pinpoint its location.,although from the available data we are not able to pinpoint its location. + ? have already argued. on the basis that a re-orientation of the central jet in ~907 sseenis physically unfeasible. for the presence of (wo highanass YSOs driving the outflows.," \citet{Guzman2010ApJ} have already argued, on the basis that a re-orientation of the central jet in $\sim90$ seems physically unfeasible, for the presence of two high-mass YSOs driving the outflows." + To compute physical parameters of the molecular outflows we followed the standard formalism described in 2.. ?— and ?.. assuming that the high velocity gas is optically (hin and ils excitational state can be described by a single excitation temperature.," To compute physical parameters of the molecular outflows we followed the standard formalism described in \citet{Bourke1997ApJ}, \citet{Garden1991ApJ} and \citet{Goldsmith1999ApJ}, assuming that the high velocity gas is optically thin and its excitational state can be described by a single excitation temperature." + II (he transitions are sub-thermally excited. (hen the derived excitation temperature would correspond to a lower limit of the kinetic temperature of the outflowing gas.," If the transitions are sub-thermally excited, then the derived excitation temperature would correspond to a lower limit of the kinetic temperature of the outflowing gas." + A general discussion of the sources oferrors have been given bv ? and ?.., A general discussion of the sources oferrors have been given by \citet{Margulis1985ApJ} and \citet{Cabrit1990ApJ}. + The main sources of error arise [from the diffieultv in determining the contribution to the outflow in (he velocity range of the ambient cloud. and not knowing the flow inclination.," The main sources of error arise from the difficulty in determining the contribution to the outflow in the velocity range of the ambient cloud, and not knowing the flow inclination." + To be conservalive we adopt as velocity boundary. between the blue and red wing emission aud the ambient emission (he values of —22.0 and —3.0... respectively.," To be conservative we adopt as velocity boundary between the blue and red wing emission and the ambient emission the values of $-22.0$ and $-3.0$, respectively." +" The column density of CO molecules in the velocity range [0],02] is given by where 74, is the rotational excitation temperature in Iv. J is the rotational quantum number of the lower state. £j,=hDJ(GJ+1) is the energy of level J. where D is the rotational constant of the CO molecule."," The column density of CO molecules in the velocity range $[v1,v2]$ is given by where $T_{\rm ex}$ is the rotational excitation temperature in K, $J$ is the rotational quantum number of the lower state, $E_J=h B J (J+1)$ is the energy of level $J$, where $B$ is the rotational constant of the CO molecule." + The Irequeney of the CO(/+14./) transition, The frequency of the transition +in Orniglia.Rich&Castro(2002):Origlia(2004) and subjected to rigorous tests in our previous studies of Galactic bulge field and cluster giants (seeOriglia&Rich&Origlia2005.andreferences therein)..,"in \citet{orc02,ori04} + and subjected to rigorous tests in our previous studies of Galactic bulge field and cluster giants \citep[see][ and references therein]{ori04,rich05}." + Reference solar abundances are [rom Sauval (1998)., Reference solar abundances are from \citet{gv98}. +. In the first iteration. we estimate stellar temperature from the (J—IxX)o. colors (see Table 1)) by using an average redddening of E(D-V)—2.9 (Schultheisetal.1999). and the color-temperature transformation of Monteerilloetal.(1998) specifically. calibrated on globular cluster giants.," In the first iteration, we estimate stellar temperature from the $\rm (J-K)_0$ colors (see Table \ref{tab1}) ) by using an average redddening of E(B-V)=2.9 \citep{sch99} + and the color-temperature transformation of \citet{MFFO98} specifically calibrated on globular cluster giants." + Gravity has been estimated [rom theoretical evolutionary. tracks. according to the location of the stars on the Red Giant Branch (RGB) (seeOrigliaοἱal.referencesthereinforamoredetailed diseussion)..," Gravity has been estimated from theoretical evolutionary tracks, according to the location of the stars on the Red Giant Branch (RGB) \citep[see][and references therein for a more detailed +discussion]{ori97}." + An average value €=2.0 kin/s has been adopted for the microturbulence (seealsoOrigliaοἱal.1997)., An average value $\xi$ =2.0 km/s has been adopted for the microturbulence \citep[see also][]{ori97}. +. More stringent constraints on (he stellar parameters are obtained bv the simultaneous spectral fitting of the several CO and OLL molecular bands. which are very sensitive to temperature. gravity. ancl mieroturbulence variations (see Figs.," More stringent constraints on the stellar parameters are obtained by the simultaneous spectral fitting of the several CO and OH molecular bands, which are very sensitive to temperature, gravity and microturbulence variations (see Figs." + 6.7 of Origlia.Rich&Castro (2002))).," 6,7 of \citet{orc02}) )." + CO and ΟΙ in particular. are extremely. sensitive (o νε in the range 3500 to 4500 KX. Indeed. temperature sets the Iraction of molecular atomic carbon and oxveen.," CO and OH in particular, are extremely sensitive to $T_{eff}$ in the range 3500 to 4500 K. Indeed, temperature sets the fraction of molecular atomic carbon and oxygen." + Al temperatures 24500 Ix molecules barely survive: most of the carbon and oxygen are in atomic form and the CO and OLI spectral features become very weak., At temperatures $\ge$ 4500 K molecules barely survive; most of the carbon and oxygen are in atomic form and the CO and OH spectral features become very weak. + On the contrary. al temperatures <3500 IX most of the carbon and oxveen are in molecular form. drastically reducing the dependence of the CO and OII band strengths and equivalent widths on the temperature itself (Orieliaetal.1997).," On the contrary, at temperatures $\le$ 3500 K most of the carbon and oxygen are in molecular form, drastically reducing the dependence of the CO and OH band strengths and equivalent widths on the temperature itself \citep{ori97}." +. The final values of our best-fit abuncances together with random errors are listed in Table 1.., The final values of our best-fit abundances together with random errors are listed in Table \ref{tab1}. . + As a further check on the statistical significancee of our best-fit solution. we also compute svnthetie spectra with ATey= z200 IN. Mogg= 40.5 dex ancl Af= x0.5 kms 1. and with corresponding simultaneous variations of the C and O abundances (on average. «0.2 dex) to reproduce the depth of the molecular features.," As a further check on the statistical significance of our best-fit solution, we also compute synthetic spectra with $\rm \Delta T_{eff}=\pm$ 200 K, $\rm \Delta log~g=\pm$ 0.5 dex and $\rm \Delta \xi=\mp$ 0.5 km $^{-1}$, and with corresponding simultaneous variations of the C and O abundances (on average, $\pm$ 0.2 dex) to reproduce the depth of the molecular features." + As a figure of merit of the statistical test. we adopt the difference between the model and the observed spectrum (herealter 9).," As a figure of merit of the statistical test, we adopt the difference between the model and the observed spectrum (hereafter $\delta$ )." + In order to quantify svstematic discrepancies. this parameter is more powerful than the classical \? test. which is instead equally sensitive to ancl systematic scatlers (seealsoOriglia2003:Origlia&Rich 2004).," In order to quantify systematic discrepancies, this parameter is more powerful than the classical $\chi ^2$ test, which is instead equally sensitive to and scatters \citep[see also][]{ori03,ori04}." +. Our best fit solutions always show 9054 probability to be representative of the observed spectra. while (hose with —0.1 dex are significant at 71 6 level.," Our best fit solutions always show $>$ probability to be representative of the observed spectra, while those with $\pm$ 0.1 dex are significant at $\ge$ 1 $\sigma$ level." + Spectral fitting solutions with abundance variations of z0.2 dex. due (o possible svstematic uncertainties of £200 Ix in temperature. £0.5 dex in gravity or $0.5 km/s in microturbulence have <30% probability of being statisticallysignificant.," Spectral fitting solutions with abundance variations of $\pm$ 0.2 dex, due to possible systematic uncertainties of $\pm$ 200 K in temperature, $\pm$ 0.5 dex in gravity or $\mp$0.5 km/s in microturbulence have $<$ probability of being statisticallysignificant." + Hence. as a conservative estimate of (he systematic error in the derived best-fit abundances. due to (he residual uncertaintv in," Hence, as a conservative estimate of the systematic error in the derived best-fit abundances, due to the residual uncertainty in" +determined.,determined. + At least the brightest portion of the Red Giant Branch (RGB) is needed in order to have hits on metallicity and distance. while the detection of the Main Sequence Turn-Off (AIS--TO) is required to get the age.," At least the brightest portion of the Red Giant Branch (RGB) is needed in order to have hits on metallicity and distance, while the detection of the Main Sequence Turn-Off -TO) is required to get the age." + In this respect. detailed investigations of simple stellar populations (SSP. coeval and chemically homogeneous stellar aggregates) offer a unique opportunity to empirically calibrate suitable photometric indices and evolutionary features in terms of the overall metallicity of the svstem.," In this respect, detailed investigations of simple stellar populations (SSP, coeval and chemically homogeneous stellar aggregates) offer a unique opportunity to empirically calibrate suitable photometric indices and evolutionary features in terms of the overall metallicity of the system." + Stellar clusters represent (hie best approximation of SSP known in the Universe. hence they are ideal tools for this purpose.," Stellar clusters represent the best approximation of SSP known in the Universe, hence they are ideal tools for this purpose." + An empirical method to get simultaneously metallicity and reddenine from the morphology and location of (he RGB in the (I.V2I) color-magnitude diagram (CMD) has been presented a decade ago (Sarajedini1994).," An empirical method to get simultaneously metallicity and reddening from the morphology and location of the RGB in the $\rm (I,V-I)$ color-magnitude diagram (CMD) has been presented a decade ago \citep{SAR94}." +. More recently. this method has been extended to the (V.B—V) plane by Sarajedini&Lavden(1997) and [further improved by Carrvetta& Dragaelia(1998). and Ferraroetal.(1999). by. adopting the Carretta&Gratton(1997). and the global ([M/II]) metallicity scale. respectively.," More recently, this method has been extended to the $\rm (V, B-V)$ plane by \citet{sarlay97} and further improved by \citet{CB98} and \citet{F99} by adopting the \citet{cg97} + and the global $\rm [M/H]$ ) metallicity scale, respectively." + The method presented here adopts the most recent calibrations as obtained [rom a laree Ih photometric database of GCs. collected by our eroup with different instruments at the ESO telescopes (La Silla. Chile) and at the TNG (La Palma. Spain) over the last decade in the framework of a long-term project devoted to study the photometric properties of the RGB (Ferraroetal.1994.1995:MontegrilloFerraro2000:Sollimaal.2004:ValentielValenti.Ferraro&Orielia 2004a.," The method presented here adopts the most recent calibrations as obtained from a large IR photometric database of GCs, collected by our group with different instruments at the ESO telescopes (La Silla, Chile) and at the TNG (La Palma, Spain) over the last decade in the framework of a long-term project devoted to study the photometric properties of the RGB \citep{F94,F95,M95,F00,S04,V04,Va04,Vb04}." +b).. In particular. we have recently presented a complete set of photometric indices (colors. magnitudes aud slopes) describing the location and the morphology of the RGB and (their calibrations in terms of the global cluster metallicitw (Valenti.Ferraro&Origlia2004a).. together with the empirical calibrations of the IR. luminosity of the two major RGB evolutionary features. namely the RGB Dump and Tip (Valenti.Ferraro&Origlia2004b).," In particular, we have recently presented a complete set of photometric indices (colors, magnitudes and slopes) describing the location and the morphology of the RGB and their calibrations in terms of the global cluster metallicity \citep{Va04}, together with the empirical calibrations of the IR luminosity of the two major RGB evolutionary features, namely the RGB Bump and Tip \citep{Vb04}." +. It is worth mentioning that. the calibration relations used in this study rely on a database. whose properties have been derived in a fully self.consistent war.," It is worth mentioning that, the calibration relations used in this study rely on a database, whose properties have been derived in a fully self–consistent way." + In fact. the adopted estimates of the cluster recldenine. metalliditv ancl distance are based on a homogeneous photometric svstemi: on the Ferraroetal.(1999). distance scale which relies on the most recent ancl largest database of Qalactie GCs: and on a uniform and highresolution metallicity scale 1991).," In fact, the adopted estimates of the cluster reddening, metallicity and distance are based on a homogeneous photometric system; on the \citet{F99} distance scale which relies on the most recent and largest database of Galactic GCs; and on a uniform and high–resolution metallicity scale \citep{cg97}." +. The calibration of suitable relations to derive metallicity. reddening and distance in the near IR. plane is crucial in the study. of extragalactic bulges. which can be characterized bv high metallicities and can be affected. by severe reddening.," The calibration of suitable relations to derive metallicity, reddening and distance in the near IR plane is crucial in the study of extragalactic bulges, which can be characterized by high metallicities and can be affected by severe reddening." + The current generation of eround-basecl IR. instrumentation with high resolution and wide field coverage. and the future availability of the James Webb Space Telescope will allow us to resolve the brightest eiants in galaxies up to several Mpc away. and (o derive their overall metallicity. reddening ancl distance modulus with great accuracy.," The current generation of ground-based IR instrumentation with high resolution and wide field coverage, and the future availability of the James Webb Space Telescope will allow us to resolve the brightest giants in galaxies up to several Mpc away, and to derive their overall metallicity, reddening and distance modulus with great accuracy." +we estimate an average VSG contamination of less than 10%.,we estimate an average VSG contamination of less than $\%$. + This implies degrading the original resolution of the Hi-GAL data by adopting the IRAS angular resolution of 4'., This implies degrading the original resolution of the Hi-GAL data by adopting the IRAS angular resolution of $^{\prime}$. +" We note, as a consequence, an averaging effect within the resolution element, which also needs to be taken into account in the interpretation of the results of the present analysis."," We note, as a consequence, an averaging effect within the resolution element, which also needs to be taken into account in the interpretation of the results of the present analysis." + One of the key aspects in the analysis of CMB is the ability to separate its emission from the other astrophysical foregrounds (including thermal dust) through multifrequency observations., One of the key aspects in the analysis of CMB is the ability to separate its emission from the other astrophysical foregrounds (including thermal dust) through multifrequency observations. +" While it is not easy to achieve successful component separation over the Galactic plane, determining the dust spectral index variations across the sky can efficiently help in reducing the number of unknowns in the problem (Ricciardietal., 2010)."," While it is not easy to achieve successful component separation over the Galactic plane, determining the dust spectral index variations across the sky can efficiently help in reducing the number of unknowns in the problem \citep{ricciardi10}." +". We use the ROMAGAL PACS and SPIRE maps described in Traficanteetal. (2010), combined with the IRIS (ImprovedReprocessingoftheIRASSurvey,seeMiville-Deschénes&Lagache,2005) 100 um data."," We use the ROMAGAL PACS and SPIRE maps described in \citet{Traficante10}, , combined with the IRIS \citep[Improved Reprocessing +of the IRAS Survey, see][]{MamD05} 100 $\mic$ data." +" Multiplicative factors (0.78,1.02, 1.05 and 0.94 at 160, 250, 350 and 500 um, respectively) have been applied to the data (Poglitschetal.,2010;Griffinetal.,Swinyard 2010)."," Multiplicative factors (0.78,1.02, 1.05 and 0.94 at 160, 250, 350 and 500 $\mic$, respectively) have been applied to the data \citep{Poglitsch10,Griffin10,Swinyard10}." +. An absolute calibration accuracy of 2096 and 1596 for PACS and SPIRE has been adopted., An absolute calibration accuracy of $\%$ and $\%$ for PACS and SPIRE has been adopted. + We applied the offsets given in Bernardet !., We applied the offsets given in \citet{Bernard10} . +". All maps have been convolved with a Gaussian kernel, with a FWHM of 4', i.e. equal to the IRIS 100 um angular resolution."," All maps have been convolved with a Gaussian kernel, with a FWHM of $^{\prime}$, i.e. equal to the IRIS 100 $\mic$ angular resolution." +" In addition, the maps have been rebinned on a uniform grid with a pixel size of 1.65'."," In addition, the maps have been rebinned on a uniform grid with a pixel size of $^{\prime}$." +" The IRIS 100 pm calibration uncertainty is taken to be 13.5% Deschénes&Lagache, 2005).", The IRIS 100 $\mic$ calibration uncertainty is taken to be $\%$ \citep[see][]{MamD05}. +. We use two different methods for the simultaneous derivation of the dust temperature and the emissivity spectral index., We use two different methods for the simultaneous derivation of the dust temperature and the emissivity spectral index. + Both allow us to fit the data with a modified blackbody function (see Eq. 1))., Both allow us to fit the data with a modified blackbody function (see Eq. \ref{eq:Inu}) ). +" In the first method we perform an y? minimization, applying the same code as described in Dupacetal.(2001)., using the IDL least-square fit function to deduce the Ty and β parameters, as well as their 1σ uncertainties."," In the first method we perform an $\chi^2$ minimization, applying the same code as described in \citet{Dupac01}, using the IDL least-square fit function to deduce the $\rm _d$ and $\beta$ parameters, as well as their $\sigma$ uncertainties." +" The second method estimates the best-fit parameters by looking for the maximum likelihood, using a Monte Carlo Markov Chain (MCMC) algorithm (Lewis&Bri- and represents a bayesian approach to parameter estimation."," The second method estimates the best-fit parameters by looking for the maximum likelihood, using a Monte Carlo Markov Chain (MCMC) algorithm \citep{Lewis02} and represents a bayesian approach to parameter estimation." +" The posterior distribution for the parameters is sampled by using the Metropolis-Hastings algorithm, and a maximum likelihood estimate is derived jointly for 6 and T4."," The posterior distribution for the parameters is sampled by using the Metropolis-Hastings algorithm, and a maximum likelihood estimate is derived jointly for $\beta$ and $T_d$." +" We chose a wide flat a priori probablity density of the parameters, i.e. 0 K<Τα «60 K and -1«f «5, in order not to constrain the fit results."," We chose a wide flat a priori probablity density of the parameters, i.e. 0 $\rm 1.0.," The amplitude criterion for large amplitude variables applied in the LMC, SMC and 6522 is $\Delta r_{\rm MACHO} > 1.0$ , whereas that for the local sample is $\Delta V > 1.0$ ." + Since» the V.. amplitudes. are usually somewhat greater than those ‘ab rstycHo. lb is !probable‘ that‘ more local‘ stars‘ have‘ been classified as having large amplitudes than should have been.," Since the $V$ amplitudes are usually somewhat greater than those at $r_{\rm MACHO}$, it is probable that more local stars have been classified as having large amplitudes than should have been." + If AV&1.63 is adopted as the criterion. the only. amplitude stars are those in or close to box C. There are several clear trends between the first. three samples.," If $\Delta V > 1.63$ is adopted as the criterion, the only large-amplitude stars are those in or close to box C. There are several clear trends between the first three samples." + Phe A. Bo and upper € sequences are relatively more populated in the Magellanic Clouds than in the Bulge.," The $^+$, $^+$ and upper $'$ sequences are relatively more populated in the Magellanic Clouds than in the Bulge." +" ""Phe low ends of the C (Mira) and € sequences are relatively uncler-populated in the Magellanie Clouds.", The low ends of the C (Mira) and $'$ sequences are relatively under-populated in the Magellanic Clouds. + Doubly periodic variables are sparse or lacking at the lower luminosity end . .. ⋜↧↓≻↓≻↓↓∢⊾∠⇂↓⊔↿↓↕∢⊾∟↳∖, Doubly periodic variables are sparse or lacking at the lower luminosity end in the Magellanic Clouds. +↓≺⊳⊾∖↓≺⋜⋯∠⊔∖≺∣≺≼≻≼≻⋅≻−≽−≽↓⊳∖∆∣⋮↼∖↓⇀∖⋖⊔∪↙⇁: . amplitude. variables. in. the €: box declinesnu with .increasing metallicity. »LMOC NGC 66522).," Further, the number of small amplitude variables in the C box declines with increasing metallicity $\rightarrow$ $\rightarrow$ 6522)." + As mentioned.. Matsunaga et al (2006). have presented a preliminary A. log/? diagram (their Fig 1) for galactic globular clusters (GC's).," As mentioned, Matsunaga et al (2006) have presented a preliminary $K$, $P$ diagram (their Fig 1) for galactic globular clusters (GCs)." + This may be compared with Fig 4.., This may be compared with Fig \ref{comparison}. + No members of the A sequence have vet been found among the GC's. though this may be a consequenceof their smaller amplitude and the fact that the survey was carried out in at," No members of the A sequence have yet been found among the GCs, though this may be a consequenceof their smaller amplitude and the fact that the survey was carried out in at" +Cepheid pulsation with current convection models.,Cepheid pulsation with current convection models. + The release of OGLE-III data. which contain a significant number of double-overtone Cepheids (Soszynsski et al. 2008))," The release of OGLE-III data, which contain a significant number of double-overtone Cepheids (Soszyńsski et al. \cite{so08b}) )" + provided additional motivation. and focused our attention on modelling the LMC pulsators.," provided additional motivation, and focused our attention on modelling the LMC pulsators." + The modelling of double-overtone pulsation and. more generally. pulsation involving higher order overtones. is not an easy task.," The modelling of double-overtone pulsation and, more generally, pulsation involving higher order overtones, is not an easy task." + The higher the order of the pulsation mode. the deeper into the envelope it penetrates.," The higher the order of the pulsation mode, the deeper into the envelope it penetrates." + Hence. to obtain reliable periods and period ratios. model envelopes should be deep and computed with higher resolution in the internal layers compared to fundamental mode Cepheic models.," Hence, to obtain reliable periods and period ratios, model envelopes should be deep and computed with higher resolution in the internal layers compared to fundamental mode Cepheid models." + On the other hand. double-overtone Cepheids are low mass objects. characterised by small growth rates.," On the other hand, double-overtone Cepheids are low mass objects, characterised by small growth rates." + This makes non-linear computations extremely time-consuming and therefore a relatively coarse mesh Is necessary to conduct an extended non-linear model survey., This makes non-linear computations extremely time-consuming and therefore a relatively coarse mesh is necessary to conduct an extended non-linear model survey. + In order to get reliable periods and period ratios with a relatively coarse mesh. the model structure. zoning. and depth of the envelope in our code were chosen to reproduce the results obtained with the LNA code of Dziembowski (1977)) as closely as possible.," In order to get reliable periods and period ratios with a relatively coarse mesh, the model structure, zoning, and depth of the envelope in our code were chosen to reproduce the results obtained with the LNA code of Dziembowski \cite{wd77}) ) as closely as possible." + Dziembowski’s code. which is coupled with the Warsaw-New Jersey stellar evolution code (see e.g. Pamyatnykh 1999)). allows the computation of deep. evolutionary models with high spatial resolution. providing accurate periods and period ratios that can be used in asteroseismic modelling (e.g. Aoskalik Dziembowski 2005).," Dziembowski's code, which is coupled with the Warsaw-New Jersey stellar evolution code (see e.g. Pamyatnykh \cite{ap99}) ), allows the computation of deep, evolutionary models with high spatial resolution, providing accurate periods and period ratios that can be used in asteroseismic modelling (e.g. Moskalik Dziembowski \cite{md05}) )." + However. the code adopts frozen- MLT convection. which is much simpler than the time-dependent treatment used in our code.," However, the code adopts frozen-in MLT convection, which is much simpler than the time-dependent treatment used in our code." + Therefore. à comparison of the model periods and period ratios computed with both codes was conducted either for purely radiative models. or with convection in the frozen-in approximation in both of the codes.," Therefore, a comparison of the model periods and period ratios computed with both codes was conducted either for purely radiative models, or with convection in the frozen-in approximation in both of the codes." + We note that with appropriate convective parameters. the model used in our envelope code can be reduced to a standard MLT (Wuchterl Feuchtinger 1998)).," We note that with appropriate convective parameters, the model used in our envelope code can be reduced to a standard MLT (Wuchterl Feuchtinger \cite{wf98}) )." +" The resulting mesh structure for our envelope models is the following: the models consist of 200 mass zones. of which 50 outer zones have an equal mass down to the anchor zone. in which the temperature is set to 7,=11000K."," The resulting mesh structure for our envelope models is the following: the models consist of $200$ mass zones, of which $50$ outer zones have an equal mass down to the anchor zone, in which the temperature is set to $T_{\rm a}=11000\,{\rm K}$." + The envelope extends down to temperatures of 8x10°K.," The envelope extends down to temperatures of $8\times 10^6\,{\rm K}$." + In the final linear and. non-linear model computations only one set of convective parameters was used namely set R3 of Table 1.., In the final linear and non-linear model computations only one set of convective parameters was used namely set R3 of Table \ref{convpar}. + The non-linear computation of double-overtone pulsation is extremely time-consuming and prohibits a more detailec parameter study., The non-linear computation of double-overtone pulsation is extremely time-consuming and prohibits a more detailed parameter study. + Therefore. the convective parameters should be carefully chosen. and we believe that set R3 represents a good choice.," Therefore, the convective parameters should be carefully chosen, and we believe that set R3 represents a good choice." + We note that with the convective paratneters of set R3. we were able to successfully model the over:ul properties of the radial velocity curves of the first overtore Cepheids (Baranowski et al. 2009)).," We note that with the convective parameters of set R3, we were able to successfully model the overall properties of the radial velocity curves of the first overtone Cepheids (Baranowski et al. \cite{bea09}) )." + With these convective parameters. the models of the radial velocity curves of the fuidamental mode Cepheids also agree with the observations (Smolee 2009b)).," With these convective parameters, the models of the radial velocity curves of the fundamental mode Cepheids also agree with the observations (Smolec \cite{rs09b}) )." +" It 1s also important that set R3 includes the effects of radiative cooling of the convective elements (y,#0)."," It is also important that set R3 includes the effects of radiative cooling of the convective elements $\gamma_{\rm r}\ne +0$ )." + The inclusion of this effect was necessary to reproduce the long periods of some of the first overtone pulsators (Baranowski et al. 2009))., The inclusion of this effect was necessary to reproduce the long periods of some of the first overtone pulsators (Baranowski et al. \cite{bea09}) ). + It was also claimed that the inclusion of radiative cooling was necessary to obtain the 10+20 non-linear double-mode models with the Florida-Budapest code (Buchler Kollátth 2000))., It was also claimed that the inclusion of radiative cooling was necessary to obtain the 1O+2O non-linear double-mode models with the Florida-Budapest code (Buchler Kollátth \cite{bk00}) ). + Finally. we would like to point out that the models to be discussed are hot and convection Is not expected to be very strong.," Finally, we would like to point out that the models to be discussed are hot and convection is not expected to be very strong." + Therefore. the results should 1ot be very sensitive to the exact values of the convective paraneters.," Therefore, the results should not be very sensitive to the exact values of the convective parameters." + The results of the linear model survey presented in this section. particularly the implications for stellar evolution theory. were published by Dziembowski Smolec (2009.. DS090 in the following).," The results of the linear model survey presented in this section, particularly the implications for stellar evolution theory, were published by Dziembowski Smolec \cite{ds09a}, DS09 in the following)." + In the present paper. the linear models provide a background for non-linear pulsation modelling. and are briefly summarised below.," In the present paper, the linear models provide a background for non-linear pulsation modelling, and are briefly summarised below." + The pulsation models were constructed along evolutionary tracks computed with the Warsaw-New Jersey stellar evolutionary code. which allows us to compute evolutionary phases before core helium ignition.," The pulsation models were constructed along evolutionary tracks computed with the Warsaw-New Jersey stellar evolutionary code, which allows us to compute evolutionary phases before core helium ignition." + The computations were conducted for two values of metallicity that are appropriate for the LMC. namely Z=0.006 and Z=0.008.," The computations were conducted for two values of metallicity that are appropriate for the LMC, namely $Z=0.006$ and $Z=0.008$." + For the hydrogen abundance. X=0.72 was adopted.," For the hydrogen abundance, $X=0.72$ was adopted." + The OP opacities (Seaton 2005)) and Asplund et al. (2004)), The OP opacities (Seaton \cite{sea05}) ) and Asplund et al. \cite{a04}) ) + solar mixture were used in the opacity computations., solar mixture were used in the opacity computations. + Rotation and overshooting from the convective core were neglected., Rotation and overshooting from the convective core were neglected. + The mixing length parameter in the evolutionary code was set to ag=1.5., The mixing length parameter in the evolutionary code was set to $\alpha_{\rm MLT}=1.5$. + In. addition. to pulsation models computed along evolutionary track (AlogLl=0.0 in the following). we have computed the models with artificially increased luminosity.," In addition to pulsation models computed along evolutionary track $\Delta\log L=0.0$ in the following), we have computed the models with artificially increased luminosity." + We imposed higher L values (AlogL=0.2 or AlogL= 0.4) at the bottom boundary of our models. keeping their mass fixed.," We imposed higher $L$ values $\Delta\log L=0.2$ or $\Delta\log L=0.4$ ) at the bottom boundary of our models, keeping their mass fixed." + The computed models cover the whole instability strip., The computed models cover the whole instability strip. + This luminosity increase was intended to model either (overlarge) overshooting during the main sequence evolution or the core helium burning phase (second and third crossings of the instability strip)., This luminosity increase was intended to model either (overlarge) overshooting during the main sequence evolution or the core helium burning phase (second and third crossings of the instability strip). + The domains of the simultaneous linear instability of the first and second overtone modes are plotted in the Petersen diagrams presented in Fig., The domains of the simultaneous linear instability of the first and second overtone modes are plotted in the Petersen diagrams presented in Fig. + 3 (upper panel for Z=0.006 and lower panel for Z= 0.008)., \ref{LMCmodels} (upper panel for $Z=0.006$ and lower panel for $Z=0.008$ ). + The corresponding period-lummosity (PL) diagrams were published in DSO9., The corresponding period-luminosity (PL) diagrams were published in DS09. + The linear models well reproduce the observed Petersen and PL diagrams., The linear models well reproduce the observed Petersen and PL diagrams. + The inferred masses of the double-overtone Cepheids are 3.0+0.5 M...," The inferred masses of the double-overtone Cepheids are $3.0\pm0.5\,{\rm M}_\odot$ ." + However. in order to reproduce the objects of longer periods (Pj.0.6 days). which represent the majority of the sample. a significant luminosity increase compared to the post- sequence evolutionary phase is required. indicating that these objects are in the core helium burning phase.," However, in order to reproduce the objects of longer periods $P_1>0.6\,{\rm days}$ ), which represent the majority of the sample, a significant luminosity increase compared to the post-main sequence evolutionary phase is required, indicating that these objects are in the core helium burning phase." + As noted by DSO9. this represents a challenge for the stellar evolution," As noted by DS09, this represents a challenge for the stellar evolution" +"of Wso using Rankin's Equation (1)), with a and P values generated by Monte Carlo simulations.","of $W_{50}$ using Rankin's Equation \ref{eq.1}) ), with $\alpha$ and $P$ values generated by Monte Carlo simulations." + Figure 6 presents comparison of Wso values obtained by these two different methods., Figure \ref{figure.6} presents comparison of $W_{50}$ values obtained by these two different methods. +" Six different ranges of o values are marked by different colours, while values of @ are represented by the length of vertical bars with the same scale as that of the ordinate axis."," Six different ranges of $\alpha$ values are marked by different colours, while values of $\beta$ are represented by the length of vertical bars with the same scale as that of the ordinate axis." +" Of course, the correlation is perfect for very small 8~0°."," Of course, the correlation is perfect for very small $\beta\sim0^{\circ}$." +" As one can see the correlation is also very good for larger 6 (corresponding to longer vertical bars), provided that inclination angles are not too small (a> 20°)."," As one can see the correlation is also very good for larger $\beta$ (corresponding to longer vertical bars), provided that inclination angles are not too small $\alpha > 20^{\circ}$ )." +" The points above and below the perfect correlation line correspond to positive and negative values of the impact angle £, respectively."," The points above and below the perfect correlation line correspond to positive and negative values of the impact angle $\beta$, respectively." +" Generally one can say, that except of cases of almost aligned rotators both methods give very similar results."," Generally one can say, that except of cases of almost aligned rotators both methods give very similar results." + This further confirms usefulness of Rankin's method expressed by her simple Equation (1))., This further confirms usefulness of Rankin's method expressed by her simple Equation \ref{eq.1}) ). +" Finally, we can make a direct comparison of Rankin's method of inclination angle estimations represented by Equation (1)) with strict calculations based on Equations"," Finally, we can make a direct comparison of Rankin's method of inclination angle estimations represented by Equation \ref{eq.1}) ) with strict calculations based on Equations" +In the case of electron (leptonic) model of Suetal.(2010). eamma-ravs are produced by scattering of relativistic electrons on background soft photons. ie. relic. IR and optical photons from the disk.,"In the case of electron (leptonic) model of \citet{meng} + gamma-rays are produced by scattering of relativistic electrons on background soft photons, i.e. relic, IR and optical photons from the disk." + In (his section we present our ideas about the origin of the Fermi Bubble in the framework of star capture by the central SMDBII., In this section we present our ideas about the origin of the Fermi Bubble in the framework of star capture by the central SMBH. + The process of eunma-ray emission from the bubble is determined by à number of stages of energy. transformation., The process of gamma-ray emission from the bubble is determined by a number of stages of energy transformation. + Each of these stages actually involves complicated plivsical processes., Each of these stages actually involves complicated physical processes. + The exact details of (hese processes are still not understood. very. well., The exact details of these processes are still not understood very well. + Nevertheless. (heir qualitative features do not depend on these details.," Nevertheless, their qualitative features do not depend on these details." + In the following. we only briefly describe these processes and give (heir qualitative interpretations.," In the following, we only briefly describe these processes and give their qualitative interpretations." + We begin to describe processes of star capture by (he central black hole as presentedin Doeieletal. (2009d).., We begin to describe processes of star capture by the central black hole as presentedin \citet{dog_aa}. . +"its abundance can then be taken as represcutative of the cosmological primordial ΠΕ,",its abundance can then be taken as representative of the cosmological primordial lithium. + The convolution of the spectra by a narrow gaussiai profile of standard deviation 1/2 pixel bringing au unsieuificant degradation in spectral resolution. but jucreasing the signal/noise ratio bv 57 per cent. make the classicalB 47Di test inapplicable.B because the noiseB on consecutive pixels becomes correlated.," The convolution of the spectra by a narrow gaussian profile of standard deviation $1/\sqrt{2}$ pixel, bringing an unsignificant degradation in spectral resolution, but increasing the signal/noise ratio by 57 per cent, make the classical $\chi^2 $ test inapplicable, because the noise on consecutive pixels becomes correlated." + In order to properly discuss the (ο.ο) residuals. i Is necessary to derive the expected mean value of the stim of the (0Coo;TOtD and its expected variance. in our specific case. where the independent random variables are the noise on each data point. couvolution.," In order to properly discuss the $(O-C)$ residuals, it is necessary to derive the expected mean value of the sum of the $((O-C)_i/\sigma_i)^2$, and its expected variance, in our specific case, where the independent random variables are the noise on each data point, convolution." + The convolved sigual 3; is expressed by a linear form of the unconvolved signal y; is the ruuniug πανο of the pixel along the Li feature): with. in our particular case. the nuuerical values: 0.01. 0.208. 0.561. 0.208. 0.0L for dw9 to uy respectively.," The convolved signal $Y_i$ is expressed by a linear form of the unconvolved signal $y_i$ is the running number of the pixel along the Li feature): with, in our particular case, the numerical values: 0.01, 0.208, 0.564, 0.208, 0.01 for $w_{-2 }$ to $w_{2}$ , respectively." + The pseudo-\?. let us call it X7: can be readily computed {1 and ζω are respectively the first aud the last pixel of the lithimim feature). assu that the y; are affected by a eaussian nolse equal to the noise on each pixel value of the spectrum before convolution ( 1/650 iu fraction of the continuum level ).," The $\chi^2 $, let us call it $X^2$: can be readily computed $i_1$ and $i_2$ are respectively the first and the last pixel of the lithium feature), assuming that the $y_i$ are affected by a gaussian noise equal to the noise on each pixel value of the spectrum before convolution ( 1/650 in fraction of the continuum level )." + It is a well known result (linear combination of gaussiau variables) that the mathematical expectation of X? is reduced with respect to the mathematical expectauce of the true 4? in the ratio: This gives the value of σ΄ The computation of he complete distribution has been performed vy ruuuiue 1 000 000 realizations. aud biuuiug the resulting values of X7.," It is a well known result (linear combination of gaussian variables) that the mathematical expectation of $X^2$ is reduced with respect to the mathematical expectance of the true $\chi^2$ in the ratio: This gives the value of $\sigma^\prime$ The computation of the complete distribution has been performed by running 1 000 000 realizations, and binning the resulting values of $X^2$." + The results are eiven in te 1 of this appendix for 223 deerees of freedom., The results are given in table 1 of this appendix for 23 degrees of freedom. + This computation shows hat the variance is 2. tics the mathematical expectauce. iustead of two for a true V6D law.," This computation shows that the variance is 3.4 times the mathematical expectance, instead of two for a true $\chi^2$ law." +" So. part of the gain iu noise reduction is lost in a shallower distribution. but the gain is auvhow substantial. as already shown. in- section. L1. where the two 47D were used competitively,"," So, part of the gain in noise reduction is lost in a shallower distribution, but the gain is anyhow substantial, as already shown, in section 4.1, where the two $\chi^2$ were used competitively." +" Several simulations. doue on test cases, have confirmed the effectiveness ofnoise filtering. with the use of a modified X7 distribution."," Several simulations, done on test cases, have confirmed the effectiveness ofnoise filtering, with the use of a modified $\chi^2$ distribution." +polarization mode. using ss iutegrations.,"polarization mode, using s integrations." + The array was in its C-configuration. sufficicutly extended to eusure that confusion did not prevent us reaching the expected thermal noise-Iinited seusitivity levels.," The array was in its C-configuration, sufficiently extended to ensure that confusion did not prevent us reaching the expected thermal noise-limited sensitivity levels." + Data reduction was performed using the Como Astronomy Software Applications (CASA) package (MeMullin.etal.2007)., Data reduction was performed using the Common Astronomy Software Applications (CASA) package \citep{McM07}. +.. Bandpass aud flux density calibration were performed using the primary flux calibrator (JOL3713309 for GRO J0122]32. aud J133113030= for NTE J1118|bso and GRO J1655-10).," Bandpass and flux density calibration were performed using the primary flux calibrator (J0137+3309 for GRO J0422+32, and J1331+3030 for XTE J1118+480 and GRO J1655-40)." + The amplitude scale was set according to the Perley-Butler coefficieuts derived from recent measurements at the EVLA., The amplitude scale was set according to the Perley-Butler coefficients derived from recent measurements at the EVLA. + Amplitude aud phase eius were derived for all calibrator sources. aud the plase calibrators used are shown in Table 1..," Amplitude and phase gains were derived for all calibrator sources, and the phase calibrators used are shown in Table \ref{tab:vla_obs}." + Finally. the calibration was applied to the target sources. which. following frequency-averagine by a factor of & were then subjected to several rounds of Huaeing and phase-ouly sef-calibration.," Finally, the calibration was applied to the target sources, which, following frequency-averaging by a factor of 8, were then subjected to several rounds of imaging and phase-only self-calibration." + To successfully nuage and deconvolve sucji wide bandwidths. we used the CASA implementation of the alegorithuu for inulti-frequency svuthesis imaging (Sault&Wicringa1991).," To successfully image and deconvolve such wide bandwidths, we used the CASA implementation of the algorithm for multi-frequency synthesis imaging \citep{Sau94}." +. Iu our original VLA survey. only Swift J1753.5-0127 was detected. with a flux density of nuuJv in each of the observations (heexactflux«cusiticshavealready.beenreportedbySolerietal. 2010).," In our original VLA survey, only Swift J1753.5-0127 was detected, with a flux density of mJy in each of the observations \citep[the exact flux densities have already been reported by][]{Sol10}." +. The other fen sources were not detected. herefore too faint for an astrometric TSA observiue campaign.," The other ten sources were not detected, and are therefore too faint for an astrometric HSA observing campaign." + Of the three svstenis that we observed witi the EVLA (GRO J0122|32. NTE J1IIS|£80. and GRO J1655-10). none were detected at a statistically siguficaut level.," Of the three systems that we observed with the EVLA (GRO J0422+32, XTE J1118+480, and GRO J1655-40), none were detected at a statistically significant level." + The observations of GRO J1655-10. were hampered by the array coufiguration (€ configuration: even the low declination of the source. a livbrid array wit ran exteuded," The observations of GRO J1655-40 were hampered by the array configuration (C configuration; given the low declination of the source, a hybrid array with an extended" +energies can onlv be obtaimed if the second absorbing uedciun is close to tjo N-raw source and has a sinular size.,energies can only be obtained if the second absorbing medium is close to the X-ray source and has a similar size. + A iaost likely candidatejs the BLR., A most likely candidate is the BLR. + Generallv. the cohunn deusitv of the partial covering medina is of he order of⋟↽⋅≻⊳↜ —107*du. 1 aereenient with estimates for the Nyy of broad line clouds.," Generally, the column density of the partial covering medium is of the order of $\rm \sim 10^{23} cm^{-2}$, in agreement with estimates for the $_H$ of broad line clouds." + The covering actor of the partial absorber is geucrally found to be larger than ~Cs that is sienificantlv lugher than the covering factor of the broad line clous which. based on the broad hues equivalent width aud ou the abseuce of the Ly-ecdee cutoff in the UV s)octruni of QSOs. is expected to be about1054.," The covering factor of the partial absorber is generally found to be larger than $\sim$, that is significantly higher than the covering factor of the broad line clouds which, based on the broad lines equivalent width and on the absence of the Ly-edge cutoff in the UV spectrum of QSOs, is expected to be about." +.. IHTowever. I «o not consider this a uajor caveat since the dual absorbers discovered so far are not represcutative of the true distribution of the partial covering absorption svstemis. but onlv sample the tai with hieh coveriug factor.," However, I do not consider this a major caveat since the dual absorbers discovered so far are not representative of the true distribution of the partial covering absorption systems, but only sample the tail with high covering factor." + Partial covering svsteus with a covering factor of probably remained uudoetected mn he past observaions., Partial covering systems with a covering factor of probably remained undetected in the past observations. + E expect Chandra and NAIN to discover a large number of dual absorbers witl1 low partial covering., I expect Chandra and XMM to discover a large number of dual absorbers with low partial covering. + Finally. thauks to the extended spectral coverage of BeppoSAX some cases of dual absorber with partial coveri1 characterize by an absorbing column as high as Ny-~--1075057 are being. foi.iid (eg.," Finally, thanks to the extended spectral coverage of BeppoSAX some cases of dual absorber with partial covering characterized by an absorbing column as high as $\rm N_H \sim 10^{24} cm^{-2}$ are being found (eg." + Turner et al., Turner et al. + 2000)., 2000). + Such a high Nyy is still consistent with that expoectecL for the broad ine clouds., Such a high $_H$ is still consistent with that expected for the broad line clouds. + Indeed. the estimatec σοι of 1022.102763 for t1ο broad line ckmids is only a lower lait that is required to produce the low ionization broad emission lines (MelII. Fell. ete...).," Indeed, the estimated column of $\rm 10^{22}-10^{23} cm^{-2}$ for the broad line clouds is only a lower limit that is required to produce the low ionization broad emission lines (MgII, FeII, etc...)." + Although the column density of the cold obscuring material does not depeud on the xoperties of the ACN (ee., Although the column density of the cold obscuring material does not depend on the properties of the AGN (eg. + its luminosity) it does depend on the properties of the rost galaxy., its luminosity) it does depend on the properties of the host galaxy. + 1deed. the absorbing Nyy strongly correlates with the presence of a stellar bar in the host ealaxyv (Ma10ino et al.," Indeed, the absorbing $_H$ strongly correlates with the presence of a stellar bar in the host galaxy (Maiolino et al." + 1999)., 1999). +") In particular. while nou-barred Sv2s are chiarveterzed by an aveπασο logNI22(eu23, most of the stroely xuvred Sv2 ealaxics are Compton tick."," In particular, while non-barred Sy2s are characterized by an average $\rm log~N_H \sim 22~(cm^{-2})$, most of the strongly barred Sy2 galaxies are Compton thick." + This fudine iudicates that stellar bars are effective in driving gas iuto the nucear reelol ito«j)bscure the ACN., This finding indicates that stellar bars are effective in driving gas into the nuclear region to obscure the AGN. +" This result is iu ine with other studies. as reviewOC by Salanoto (1999). which indicate that ALS are effective in driving eas mto the uuclear regiou. though these other stucies are rot specifically focused on Ανα,"," This result is in line with other studies, as reviewed by Sakamoto (1999), which indicate that bars are effective in driving gas into the nuclear region, though these other studies are not specifically focused on AGNs." + We mieh speculate that. more generally. 1011”axisviunietrie poteutials (ee.," We might speculate that, more generally, non-axisymmetric potentials (eg." + distored morphologies auc galaxy interactions) crive eas iuto the nuclear region., distorted morphologies and galaxy interactions) drive gas into the nuclear region. + Althoeh there are sole observational indications nu lis direction. a svstende study. siilar to hat on barred systems. has not been vertormed vet.," Although there are some observational indications in this direction, a systemic study, similar to that on barred systems, has not been performed yet." + This wi Lhe possible with the new Chandra and NMM data., This will be possible with the new Chandra and XMM data. + If the obscurueg torus as the same gas-to-dust ratio as iu the Calactic ISAL. aud the dust is characterizes La Galactic extinction curve. then the unclear region of Sv2s should suffer a visual «sxtinction that is related to the gaseous column density by the formmla Ay=5x107?Ni(em7).," If the obscuring torus has the same gas-to-dust ratio as in the Galactic ISM, and the dust is characterized a Galactic extinction curve, then the nuclear region of Sy2s should suffer a visual extinction that is related to the gaseous column density by the formula $\rm A_V=5\times 10^{-22} N_H +(cm^{-2})$." + In general this is not the case: Ay is lower than expected from the Nyy measured in the N-ravs., In general this is not the case: $\rm A_V$ is lower than expected from the $_H$ measured in the X-rays. + This was first pointed out bv Maccacaro et al. (, This was first pointed out by Maccacaro et al. ( +1982).,1982). + A visual extinction lower than that expected from the Ny measured in the N-ravs is also required to fit the IR spectrmu of AGNs (Cranato, A visual extinction lower than that expected from the $_H$ measured in the X-rays is also required to fit the IR spectrum of AGNs (Granato +(Usb Bee),(t) - )^2. +", Encrey couservation for the PWN system reads: Enel?) Eat)— Eust) | Lot :", Energy conservation for the PWN system reads: (t) + (t)= (t) + L_0 t . + Tere £u) is the kinetic enerev which the swept- ejecta would have if they were freely expanding., Here $E_{\rm init}(t)$ is the kinetic energy which the swept-up ejecta would have if they were freely expanding. + This quantity cau be obtained by iutegratiug the kinetic energy deusity of ejecta in a sphere with radius ¢«aif there was no PWN., This quantity can be obtained by integrating the kinetic energy density of ejecta in a sphere with radius $r < R_{\rm pwn}$ if there was no PWN. + This vields: , This yields: (t)=E_0 )^5. +After some algebra using the equations (1))-(11)) one can obtain a power-law solution for the radius of the pulsar wind bubble:ο... where C is à nunerical coustaut of order uuitv: Revuolds aud Chevalier(198D). already: obtained this Rit)x(07 expausion law.," After some algebra using the equations \ref{Density}) \ref{Econs}) ) one can obtain a power-law solution for the radius of the pulsar wind bubble:, where $C$ is a numerical constant of order unity: Reynolds and Chevalier(1984) already obtained this $R(t)\propto t^{6/5}$ expansion law." + It can casily be checked that the expansion velocity obtained iu this manner is indecd wach larecr than the sound. velocity in the freely expanding supernova remmant., It can easily be checked that the expansion velocity obtained in this manner is indeed much larger than the sound velocity in the freely expanding supernova remnant. + Towards the cud of the free expansion stage a reverse shock is driven deep into the interior of the SNR., Towards the end of the free expansion stage a reverse shock is driven deep into the interior of the SNR. + This reverse shock reheats the stellar ejecta. aud as a result the sound velocity increases by a large factor.," This reverse shock reheats the stellar ejecta, and as a result the sound velocity increases by a large factor." + When the reverberations due to reflections of the reverse shock lave alinost completely dissipated. one ean approximate the interior of the SNR by using the analytical Sedov solution (Sedov. 1958)).," When the reverberations due to reflections of the reverse shock have almost completely dissipated, one can approximate the interior of the SNR by using the analytical Sedov solution \cite{Sedov}) )." + The interaction with the reverse shock iufluences the evolution of the pulsar wind uchula quite dramatically., The interaction with the reverse shock influences the evolution of the pulsar wind nebula quite dramatically. + Ciofh et ((1988) have already shown iu them 1D simulation of a pure shell SNR that the reverse shock gives rise o all kinds of sound waves and weak shocks traveling ck and forth through the ejecta before hne iuterior relaxes towards a Sedov solution., Cioffi et (1988) have already shown in their 1D simulation of a pure shell SNR that the reverse shock gives rise to all kinds of sound waves and weak shocks traveling back and forth through the ejecta before the interior relaxes towards a Sedov solution. + We will show that during he process of relaxation the radius of the oulsar winel jiebula contracts aud expands due to reverberatious of he reverse shock., We will show that during the process of relaxation the radius of the pulsar wind nebula contracts and expands due to reverberations of the reverse shock. + Conrpression waves are partlv reflected and partly transmitted at the οσο of the PWN., Compression waves are partly reflected and partly transmitted at the edge of the PWN. + We will come back to this point when we discuss results from wadrodyvuaiics simulations i section 1. which allow a nore detailed picture of this process.," We will come back to this point when we discuss results from hydrodynamics simulations in section 4, which allow a more detailed picture of this process." + Iu this Section we cousicer a fully relaxed Sedov SNR., In this Section we consider a fully relaxed Sedov SNR. + The PAWN expands subsonically into the roiinant because he interior of the SNR las been re-heated by the reverse shock., The PWN expands subsonically into the remnant because the interior of the SNR has been re-heated by the reverse shock. + For the case of a coustaut (amechanical) huuinosity daiving the pulsar wind an analytical expression for the radius of the PWN can be easily obtained., For the case of a constant (mechanical) luminosity driving the pulsar wind an analytical expression for the radius of the PWN can be easily obtained. + Iu this stage of the PWN evolution. we associate its radius Zia with the coutact discoutiuuitv separating pulsar windmaterial from the ejecta of the progenitor star (sec figure 2).," In this stage of the PWN evolution, we associate its radius $R_{\rm pwn}$ with the contact discontinuity separating pulsar windmaterial from the ejecta of the progenitor star (see figure 2)." +Iu this section. we show how the risk function IU) releq:riskhat)) cai be iminunized subject to the monotonicity constraiuts l1> 0.,"In this section, we show how the risk function $\widehat{R}(\lambda)$ \\ref{eq:riskhat}) ) can be minimized subject to the monotonicity constraints $1 \ge \lambda_0 \ge \lambda_1 \ge \ldots \ge \lambda_{N-1} \ge 0$ ." + The risk fiction correspouding to the unweighted loss function (M=ta) has a simple welelitec-stuir-of-squares form. aud cau be minimized exactly aud efficiently. using the pooled adjacent violators (PAV) algorithm (?)..," The risk function corresponding to the unweighted loss function $W = I_N$ ) has a simple weighted-sum-of-squares form, and can be minimized exactly and efficiently using the pooled adjacent violators (PAV) algorithm \citep{RWD1988}." + While the risk function corresponding to the loss fueion (WW4 £s) is still quadratic in A. it can no louger be expressed as a weighted sunmr-ol-seuares. ald the PAV aleorithim cannot be used to minimize it.," While the risk function corresponding to the inverse-noise-weighted loss function $W \ne I_N$ ) is still quadratic in $\lambda$, it can no longer be expressed as a weighted sum-of-squares, and the PAV algorithm cannot be used to minimize it." +" It can be shown that. clisregardiug terms that do not depend ou A. the risk functiou (A) relec:riskhat)) can be written as where Hj,=25Muhi(H—V).1...1)7. and Vj,=2Wj.Bjj."," It can be shown that, disregarding terms that do not depend on $\lambda$ , the risk function $\widehat{R}(\lambda)$ \\ref{eq:riskhat}) ) can be written as where $H_{jk} = 2 z_j z_k W_{jk}, h = ( H - V ) (1,1,\ldots,1)^T,$ and $V_{jk} = 2 W_{jk} B_{kj}$." + H and V are both uauifestly sviumetric., $H$ and $V$ are both manifestly symmetric. + Positive (semi)defiuiteuess of W implies that Η is a positive (semi)delinite ualrix. iuplyiug that (A) is a convex function.," Positive (semi)definiteness of $W$ implies that $H$ is a positive (semi)definite matrix, implying that $\widehat{R}(\lambda)$ is a convex function." + The system of linear inequality coustraints implies that the coustrained region 'effigmonotoue)) has a convex trianguloidal shape determined by flat surfaces., The system of linear inequality constraints $1 \ge \lambda_0 \ge \lambda_1 \ge \ldots \ge \lambda_{N-1} \ge 0$ implies that the constrained region \\ref{fig:monotone}) ) has a convex trianguloidal shape determined by flat surfaces. + The original risk uiuimization problem can thereforebe formulated. as the following equivalent convex. quadratic ninimization problem: where C' is the CN—1)xN matrix, The original risk minimization problem can thereforebe formulated as the following equivalent convex quadratic minimization problem: where $C$ is the $(N-1) \times N$ matrix -1 1 +These show no evidence for any dependence on pulsation period.,These show no evidence for any dependence on pulsation period. +l. Therefore we take a weighted mean for each ratio: the distance calculations differ by 1.5%40.6%. with the Davesian being larger. and the radius calculations differ by 1.12:0.776. with the Davesian again being larger.," Therefore we take a weighted mean for each ratio: the distance calculations differ by $1.5\%\pm0.6\%$, with the Bayesian being larger, and the radius calculations differ by $1.1\%\pm0.7\%$, with the Bayesian again being larger." + For comparison. the best individual distance and radius measurements in our dataset are for X Cvg. for which the Bavesian uncertainties are both 42.6%. larger than the possible svstemiatic difference between (he Davesian aud bisector calculations.," For comparison, the best individual distance and radius measurements in our dataset are for X Cyg, for which the Bayesian uncertainties are both $\pm2.6\%$, larger than the possible systematic difference between the Bayesian and bisector calculations." + Our first result is that the distances ancl radii computed by the (wo methods agree quite well., Our first result is that the distances and radii computed by the two methods agree quite well. + In Figures 4 and 5 (he uncertainties can be seen (o increase as (he period is shorter., In Figures \ref{distratio} and \ref{radratio} the uncertainties can be seen to increase as the period is shorter. + This is likely à result of the smaller pulsation amplitudes at shorter periods. which result in the photometric aud velocity uncertainties having greater οδοί on the computed distances and racii.," This is likely a result of the smaller pulsation amplitudes at shorter periods, which result in the photometric and velocity uncertainties having greater effect on the computed distances and radii." + More importantly. a glance al Table 1. or Figures 4 ancl 5.. shows that the uncertainties in (he distances and radii disagree substantially between the (wo calculations.," More importantly, a glance at Table 1, or Figures \ref{distratio} and \ref{radratio}, shows that the uncertainties in the distances and radii disagree substantially between the two calculations." + Typically the Bavesian uncertainty is more than three times tlie linear-bisector uncertainty., Typically the Bayesian uncertainty is more than three times the linear-bisector uncertainty. + Because the concept of an uncertaintv estimated from the Bavesian AICAIC posterior probability distribution may not be clear. we show an example in Figure 6.. (he posterior probability distribution for the distance to U Ser. which star is (wpical of our results.," Because the concept of an uncertainty estimated from the Bayesian MCMC posterior probability distribution may not be clear, we show an example in Figure \ref{histogram}, the posterior probability distribution for the distance to U Sgr, which star is typical of our results." + on the probability distribution is a normal distribution constructed for the same distance (592 pc). sigma (421 pc) and area (10.000 samples).," Over-plotted on the probability distribution is a normal distribution constructed for the same distance (592 pc), sigma $\pm21$ pc) and area (10,000 samples)." + The normal distribution describes (he posterior probability distribution lor the distance verv well., The normal distribution describes the posterior probability distribution for the distance very well. + This justifies our adopling the siema of the corresponding normal distribution as a lo estimator [or ihe uncertainty in the distance (and similarly for the radius) determined in the Bavesian caleulation., This justifies our adopting the sigma of the corresponding normal distribution as a $\sigma$ estimator for the uncertainty in the distance (and similarly for the radius) determined in the Bayesian calculation. + This estimator is then compared to the lo estimator from the linear-bisector computation., This estimator is then compared to the $\sigma$ estimator from the linear-bisector computation. + We add a caveat to the previous paragraph., We add a caveat to the previous paragraph. + Because our computation determines the stellar parallax. not the stellar clistance. (see equation (10))) the posterior probability distribution for the distance can become asvimmetric when the errors are large.," Because our computation determines the stellar parallax, not the stellar distance, (see equation \ref{eq:likelihood}) )) the posterior probability distribution for the distance can become asymmetric when the errors are large." + As the uncertainties become large. (he parallax posterior probability distribution becomes broad (large sigma).," As the uncertainties become large, the parallax posterior probability distribution becomes broad (large sigma)." + Its reciprocal. (he distance posterior probability distribution. will also become broaddistances.," Its reciprocal, the distance posterior probability distribution, will also become broad." + The same asymmetry will arise for the racius posterior probability distribution in such cases because (hat distribution is the product of the angular diameter posterior probability distribution (svimmietric) and the distance, The same asymmetry will arise for the radius posterior probability distribution in such cases because that distribution is the product of the angular diameter posterior probability distribution (symmetric) and the distance +Extragalactic. radio. sources were dividedu into. two morphological. classes by :2. FRILL,Extragalactic radio sources were divided into two morphological classes by \citet{fr74}. + sources. have— an cedge-darkened structure. whereas FRILL sources are edge-brightened.: with. prominent. outer. hot-spots.," I sources have an edge-darkened structure, whereas II sources are edge-brightened with prominent outer hot-spots." + MiThis Classification⊀⋅⊀ has proved. to be extremely robust: the division between the classes depends: primarily on. radio uminosity (2).. with FRU sources being more powerful. out. also on the stellar Iuminosity of the host galaxy. (?)..," This classification has proved to be extremely robust: the division between the classes depends primarily on radio luminosity \citep{fr74}, with II sources being more powerful, but also on the stellar luminosity of the host galaxy \citep{lo96}." + sPhere are significant.⋠⋠⋅ differencess between the structures of . . uM ⇂↥⋖⋅↥⋖⊾↥⋡∖↓⊔⇂↥⋖⊾⇂∖∖⊽∪≼∙⇂∥⋡∖⋡∖⋖⊾⋡∖∶⇂↓↥∪⋡∖⋖⊾↓⊔↓⊲∐∐⋡∖∪⊔↓⋅⊓⊾⋡∖∩∐⋖⊾⊔∐⋜⊔⋅⋖⊾⋅ close to the nucleus and have large opening. angles. whereas heir. equivalents: inmM UL sources are highly: eollimated. out o the hot-spots .(2)..," There are significant differences between the structures of the jets in the two classes: those in I sources often flare close to the nucleus and have large opening angles, whereas their equivalents in II sources are highly collimated out to the hot-spots \citep{bridle84}." + sPhere is. good evidence. that 1 jets. are initially relativistic. but decelerate on kKiloparsec scales. whereas jets remain relativistic until they terminate citealt Laing03)).," There is good evidence that I jets are initially relativistic, but decelerate on kiloparsec scales, whereas II jets remain relativistic until they terminate \\citealt{Laing93}) )." + The process of deceleration in PRI jets appears to be complex. and max involve a transition to turbulent Dow.," The process of deceleration in I jets appears to be complex, and may involve a transition to turbulent flow." + In addition. the sources have a wide range of morphologies. ranging. from⋅ well-defined⋅ lobes similar⊀⊀ to those in∢∖≽ FRILL sources to extended: plumes or tailsqo 19(2)..," In addition, the sources have a wide range of morphologies, ranging from well-defined lobes similar to those in II sources to extended plumes or tails \citep*{parma96}." + For: these reasons. attempts to construct elobal models of the evolution⊀ ofDEN LL sources. linkingqs observable quantities↔ such as linear. size. ancl radio. Luminosity.⊀⊀ have been less straightforward.⋅ than the equivalents for LLL sources citealt *Scheuerrd.ka07.kda07)). which assume that the jet lows are essentially. laminar.," For these reasons, attempts to construct global models of the evolution of I sources, linking observable quantities such as linear size and radio luminosity, have been less straightforward than the equivalents for II sources \\citealt*{Scheuer74,ka97,kda97}) ), which assume that the jet flows are essentially laminar." +. κPart of⋅ the motivation⊀⊀ for⋅ the . ⋅ ↓≻↓⋅⋖⋅≱∖∢⋅⊔↥≱∖↿⋯⇂∙∖⇁↓≱∖↥∐⋖⋅↓⋅∢⋅⇂∪↓⋅⋖⋅↿∪≼∙∪⊔⊳∖↿↓⋅⋯⇍↿⋜↧⊳∖↓⊔↓↓≻↓⋖⊾⊔↓⋯⇂⋖⊾↓∪⇂ . FRAIL ∖≽⊲jets for⋅ use as input. to global models., Part of the motivation for the present study is therefore to construct a simple model of I jets for use as input to global models. + We] consider. twin-jetD. sources. which. make up at least one half ⋅⋅of the IL porxopulation (?).. exeludingP wide-angle2 tail ancl fat-double sources (22). whose jet properties ciller significantlv.," We consider twin-jet sources, which make up at least one half of the I population \citep{parma96}, excluding wide-angle tail and fat-double sources \citep{ol89,ow91}, whose jet properties differ significantly." + Over the last low vears. detailed moclelling of deep VLA observations of jets in five FRE sources has allowed. us to quantify their geometries. velocity. distributions. magnetic fielcls and emissivity cistributions in three dimensions.," Over the last few years, detailed modelling of deep VLA observations of jets in five I sources has allowed us to quantify their geometries, velocity distributions, magnetic fields and emissivity distributions in three dimensions." + We refer in detail to the analvsis of by 2.hereafter LBO2a:: observations and models of a 331.further four sources have subsequently.been published. (222)...," We refer in detail to the analysis of 31 by \citet[hereafter +LB02a]{lb02a}; ; observations and models of a further four sources have subsequentlybeen published \citep{cl04,canvin05,lcbh06}. ." + NX consistent, A consistent +"The upper and lower mass-euts. z;j, and 7. respectively. represents the high-mass trunction and the brown-dwarf limit.","The upper and lower mass-cuts, $m_{\rm u}$ and $m_{\rm l}$, respectively, represents the high-mass trunction and the brown-dwarf limit." +" In the IRA the gas-return fraction A of a generation of stars is obtained from the convolution of the IMF and the stellar remnant mass wii). Le. The remnant mass wz) is a function of the initial stellar mass m. which may be approximated by a simple power-law. km"". where & is a constant."," In the IRA the gas-return fraction $R$ of a generation of stars is obtained from the convolution of the IMF and the stellar remnant mass $w(m)$, i.e., The remnant mass $w(m)$ is a function of the initial stellar mass $m$, which may be approximated by a simple power-law, $w(m) = k\, m^n$ , where $k$ is a constant." + Using the remnant masses obtained by vandenHoek&Groenewegen(1997) and Woosley&Weaver(1995) for LIM and stars. respectively.one obtains Kk=O.5 and n=0.4 for stellar masses less than a7=35M. (see Fig. 1».," Using the remnant masses obtained by \citet{vandenHoek97} and \citet{WW95} for LIM and stars, respectively, one obtains $k=0.5$ and $n=0.4$ for stellar masses less than $m = 35 M_\odot$ (see Fig. \ref{remnants}) )." + The lock-up fraction a. used in the model equations discussed in Sect. 2.3.," The lock-up fraction $\alpha$, used in the model equations discussed in Sect. \ref{equations}," + is obtained as a=|—A., is obtained as $\alpha = 1-R$. + Adopting an IMF of the form shown above in Eq. (22)), Adopting an IMF of the form shown above in Eq. \ref{larson}) ) + and defining Gta.b.c.z)=Tta.z/c)- TMa.c/b). where Fta.z) is the incomplete gamma-function. one arrives at the expression (which holds for «> Lom< 1) where v'2x—Il. i2n— Land gj is the minimum stellar mass effectively contributing to the matter cycle at the considered redshift.," and defining $G(a,b,c,z) \equiv \Gamma(a,z/c) - \Gamma(a,z/b)$ , where $\Gamma(a,z)$ is the incomplete gamma-function, one arrives at the expression (which holds for $x>1$, $n\le 1$ ) where $x' = x-1$, $n' = n-1$ and $m_\tau$ is the minimum stellar mass effectively contributing to the matter cycle at the considered redshift." +" With the parameter values given inTable | and zn=3.0M.; (the initialmass for which the stellar lifetime corresponds to the typical age of galaxies at z=3— 6) the lock-up fraction for a normal Larson(1998) IMF ο.= 0.3547.) becomes a=0.63. and in the top-heavy Ga,= 10.047.) case a=0.10."," With the parameter values given inTable \ref{imfpar} and $m_\tau = 3.0 M_\odot$ (the initialmass for which the stellar lifetime corresponds to the typical age of galaxies at $z = 5-6$ ) the lock-up fraction for a normal \citet{Larson98} IMF $m_{\rm c} = 0.35 M_\odot$ ) becomes $\alpha = 0.63$, and in the top-heavy $m_{\rm c} = 10.0 M_\odot$ ) case $\alpha = 0.10$." + As shown in the right panel of Fig. 1. e," As shown in the right panel of Fig. \ref{remnants}," +" is a steep function of the turn-over mass m. but only weakly dependent on z7,."," $\alpha$ is a steep function of the turn-over mass $m_{\rm c}$, but only weakly dependent on $m_\tau$." + Hence. the choice of m; is not critical for either a. nor the effective primary yield vy.," Hence, the choice of $m_\tau$ is not critical for either $\alpha$, nor the effective primary yield $y_{\rm d}$." + Gall(2010). computed the total dust. productivity of a generation of stars considering ditferent IMFs and stellar production efficiencies. ejm)=vqUm/vzUnm.," \citet{Gall10thes} computed the total dust productivity of a generation of stars considering different IMFs and stellar dust-production efficiencies, $\epsilon_{\rm d}(m) \equiv y_{\rm d}(m)/y_Z(m)$." + Accordig to Gall(2010) the dust production ethciency derived from observations of SN remnants is well-approximated by a power law. which connects nicely with theoretically expected ethciencies for AGB stars (Ferrarotti&Gail2006).," Accordig to \citet{Gall10thes} the dust production efficiency derived from observations of SN remnants is well-approximated by a power law, which connects nicely with theoretically expected efficiencies for AGB stars \citep{Ferrarotti06}." +. If this is a general law for stellar dust production. SNe are insignificant dust producers when weighted by an IMF.," If this is a general law for stellar dust production, SNe are insignificant dust producers when weighted by an IMF." + Theoretical work. as well as a handfull of studies. allows for a much higher etficieney in SNe.," Theoretical work, as well as a handfull of studies, allows for a much higher efficiency in SNe." + Using the results of Todini&Ferrara(2001). Gall(2010). find that AGB stars would be insignificant relative to SNe as dust producers if no dust destruction due to the reverse shock.," Using the results of \citet{Todini01}, \citet{Gall10thes} find that AGB stars would be insignificant relative to SNe as dust producers if no dust destruction due to the reverse shock." + The integrated (IMF-weighted) yields calculated by (2010) is used in this paper. Le.. three cases are considered (see Table 23: “observed” SN dust yields. theoretical SN yields with dust destruction due to reverse shocks Chigh’yield.Bianchi&Schneider 2007).. and the theoretical “upper limit’according to Todini&Ferrara(2001.referredtoasmaximalyield) — all three in combination with the theoretical results for AGB stars (Ferrarotti&Gail2006:Zhukovska.Trieloff 2008).," The integrated (IMF-weighted) yields calculated by \citet{Gall10thes} is used in this paper, i.e., three cases are considered (see Table \ref{yields}) ): 'observed' SN dust yields, theoretical SN yields with dust destruction due to reverse shocks \citep['high' yield,][]{Bianchi07}, and the theoretical 'upper limit' to \citet[][referred to as 'maximal' yield]{Todini01} – all three in combination with the theoretical results for AGB stars \citep{Ferrarotti06,Zhukovska08}." +. A ‘closed-box” evolution is obviously incompatible with the scenario described above in Sect. 2.1.., A 'closed-box' evolution is obviously incompatible with the scenario described above in Sect. \ref{cosmology}. + But for completenes. and since it corresponds to the scenario studied by Valiante as Well as Galletal.(Q0LOb).. it is included also here.," But for completenes, and since it corresponds to the scenario studied by \citet{Valiante09} as well as \citet{Gall10b}, it is included also here." + If there is no infall. the rate of change of stellar mass is Assuming no dust destruction. combination of Eq. (169) ," If there is no infall, the rate of change of stellar mass is Assuming no dust destruction, combination of Eq. \ref{dem3}) )" +and Eq. (109) , and Eq. \ref{dem4}) ) +gives the classical elosed-box solution. Replacing vy with the metal yield vy. this is also the solution for the metallicity Z in a closed-box scenario.," gives the classical closed-box solution, Replacing $y_{\rm d}$ with the metal yield $y_Z$, this is also the solution for the metallicity $Z$ in a closed-box scenario." +" If the destruction term is included.the solution reads (Edmunds2001:Dweketal. 2007). This solution differs from the classical closed-box soultion in that Zy xy/Y as M,/M,—oo,"," If the dust-destruction term is included,the solution reads \citep{Edmunds01,Dwek07}, This solution differs from the classical closed-box soultion in that $Z_{\rm d}\to y_{\rm d}/\nu$ as $M_{\rm s}/M_{\rm g} \to \infty$." + Assuming that dust grains may form and grow out of metals produced by previous generations of stars. Le.. a kind of “secondary” dust. one may redefine the dust yield as where € is the fraction of the metalspresent in the ISM that will form dust grains.," Assuming that dust grains may form and grow out of metals produced by previous generations of stars, i.e., a kind of 'secondary' dust, one may redefine the dust yield as where $\epsilon$ is the fraction of the metalspresent in the ISM that will form dust grains." + There is a natural upper limit to the value of € since Z4€ Z. ie. available.," There is a natural upper limit to the value of $\epsilon$ since $Z_{\rm d} \le Z$ , i.e., ." + With the modification above. the solution for Zy becomes (see where w= e/v.," With the modification above, the solution for $Z_{\rm d}$ becomes \citep[see also][]{Edmunds01} + where $\omega \equiv \epsilon/\nu$ ." +" For the special case ójg,.=ϐ (no dust destruction) the solution is Assume the rate of gas consumption is exactly balanced by the rate of infall.", For the special case $\delta_{\rm ISM}=0$ (no dust destruction) the solution is Assume the rate of gas consumption is exactly balanced by the rate of infall. + With this assumption made.and no dust destruction term. the solution to Eq. (103) ," With this assumption made,and no dust destruction term, the solution to Eq. \ref{dem4}) )" +is the “extreme infall model’ by (1972). Le.," is the 'extreme infall model' by \citet{Larson72}, , i.e.," +he lines.,the lines. +" Unfortunately. the PALAS spectra show that the Πο IT A1686 /,O III]A5007 cuiission region. and the PMAS uosaic do not cover the ecutive Holubere II region 42770."," Unfortunately, the PMAS spectra show that the He II $\lambda4686$ $[$ O $]~\lambda5007$ emission region, and the PMAS mosaic do not cover the entire Holmberg II region 70." + The lower pancl of Fig., The lower panel of Fig. + |. shows that the. PALAS FWIIM of the ΤΠ) line inside the I IT region is iu geucral arger compared to that of the forbidden |O III] A5007 iuc., \ref{pmas_flux} shows that the PMAS FWHM of the $\beta$ line inside the H II region is in general larger compared to that of the forbidden [O III] $\lambda5007$ line. + The PMAS FWIIM of both lines peak around the xosition of the ACIS-S error circle., The PMAS FWHM of both lines peak around the position of the ACIS-S error circle. + Because the PATAS and AIPFS EWIIM maps are giveu with different scales we present for comparison the reas values of the FWIIA inside he II IT reeion derived from both iutrunents iu Table L., Because the PMAS and MPFS FWHM maps are given with different scales we present for comparison the peak values of the FWHM inside the H II region derived from both intruments in Table \ref{peakFWHM}. + The MPFS FWIIM reals values are in agreenient with those values derived TOlu PMAS (seeping in mind the lower angular aud spectral resolution of the MPFS data}. except for [S TI} A6TIT. which is probably bleuded with the |8 ΤΠ AG731 emission iuc.," The MPFS FWHM peak values are in agreement with those values derived from PMAS (keeping in mind the lower angular and spectral resolution of the MPFS data), except for [S II] $\lambda6717$, which is probably blended with the [S II] $\lambda6731$ emission line." + The FWHAL of I} derived frou PALAS reaches about 30 laws + inside the II IL region. but outside the Ue I sub-reegion. it is up οδ kins +.," The FWHM of $\beta$ derived from PMAS reaches about 30 km $^{-1}$ inside the H II region, but outside the He II sub-region, it is up to 80 km $^{-1}$." +" The FWIIM of the 30 au + inside the ID II region (which trauslates into a velocity dispersion of about 13 kins lj is consistent with he velocity dispersion measurements of gy,=13.02:0.5 aus + and Torydsoor=1340.2 kin ! obtained by Ilippeleiu (1986))."," The FWHM of the 30 km $^{-1}$ inside the H II region (which translates into a velocity dispersion of about 13 km $^{-1}$ ), is consistent with the velocity dispersion measurements of $\sigma_{H\alpha}=13.0\pm0.3$ km $^{-1}$ and $\sigma_{[O III] \lambda5007}=11.3\pm0.2$ km $^{-1}$ obtained by Hippelein \cite{Hip86}) )." + However. the FUIWAL more than doubles at the X- position.," However, the FHWM more than doubles at the X-ray position." + Assuming the virial theorem and that the enclosed: mass is related to the velocity chauge of ~50 lan | at the distance of ~30 pe the resultine mass of a putative black hole would be about LOAD:., Assuming the virial theorem and that the enclosed mass is related to the velocity change of $\sim$ 50 km $^{-1}$ at the distance of $\sim$ 30 pc the resulting mass of a putative black hole would be about $10^{7}$ $_{\sun}$. + Therefore we believe that the increased EWIIM at the location of the Πο IL/N-rav. source may be nuderstood as ανασα] influence of the putative black hole like the accretion disk. wind or jets., Therefore we believe that the increased FWHM at the location of the He II/X-ray source may be understood as dynamical influence of the putative black hole like the accretion disk wind or jets. + The MPFES FWIIM maps show a very complex velocity field in the South-East direction of the position. mut outside the main II IT region.," The MPFS FWHM maps show a very complex velocity field in the South-East direction of the position, but outside the main H II region." +" Furthermore there is a οσα] of ~150—200 kms ft at about 5"" South-West of he X-ray source position iu the AIPFS FWIIDAL maps of Il)Als6l. |O TH) A5007. aud Ho."," Furthermore there is a peak of $\sim150-200$ km $^{-1}$ at about $^{\prime\prime}$ South-West of the X-ray source position in the MPFS FWHM maps of $\beta~\lambda4861$, [O III] $\lambda5007$, and $\alpha$." + The complex velocity field outside the II II region is most probably not related ο Toll X-1., The complex velocity field outside the H II region is most probably not related to HoII X-1. + To determine the racial velocity feld around Toll N-1 we have used the LSS spectra at the slit positious N2 aud N3 aud the MPFS data., To determine the radial velocity field around HoII X-1 we have used the LSS spectra at the slit positions N2 and N3 and the MPFS data. + Iu Fig., In Fig. + 11 we show the coutimmun and backeround subtracted line isophotes of the 2d-LSS spectra in the TTA4686 cinission line., \ref{isoph} we show the continuum and background subtracted line isophotes of the 2d-LSS spectra in the $\lambda 4686$ emission line. + The position of the cussion line (or equivaleuth. the position of the line cimittine reeion) along the wavelength cirection (accros the slit) is eiven as an offset in kins +.," The position of the emission line (or equivalently, the position of the line emitting region) along the wavelength direction (accros the slit) is given as an offset in km $^{-1}$." +" The position along the slit is eiven as an offset in arcsec. where East edge of the hec of the compact II II region #770 is located at au offset of ~17"" aud the Πο IT region at an offset of ~11.5”."," The position along the slit is given as an offset in arcsec, where East edge of the ”heel” of the compact H II region 70 is located at an offset of $\sim17^{\prime\prime}$ and the He II region at an offset of $\sim14.5^{\prime\prime}$." + The offset values decrease in the North-West direction., The offset values decrease in the North-West direction. + Both spectra N2 and. N3 cover IToIT X-1., Both spectra N2 and N3 cover HoII X-1. + The HI peak intensity iu N2 is 1.1 times ereater than that in N23., The II peak intensity in N2 is 1.4 times greater than that in N3. + The TT region is located near the edge of the IETIT region., The II region is located near the edge of the II region. + Alone the slits we can directly study the structure of the region aud we can see the complex structure of the IIT cussion., Along the slits we can directly study the structure of the region and we can see the complex structure of the II emission. + Even the shift in the slit position from N2 to N3 (0.67) results in a notable change of the isopliote structures., Even the shift in the slit position from N2 to N3 $^{\prime\prime}$ ) results in a notable change of the isophote structures. + The line isophote structures across the slits cau not be directly interpreted as radial velocity variations. because," The line isophote structures across the slits can not be directly interpreted as radial velocity variations, because" +closely follows that described by Torresetal.(2001) and consists of calculating svuthetic light curves that result from the three objects for a wide rauge of eclipsinto binary parameters. and conrparius those light curves against the photometry for. the candidate.. inB à XD scuse.,"closely follows that described by \citet{torr2004} and consists of calculating synthetic light curves that result from the three objects for a wide range of eclipsing binary parameters, and comparing those light curves against the photometry for the candidate, in a $\chi^2$ sense." + We regard as acceptable any. blend sceuario that results iu a synthetic light curve giving a 4 for the ft that is not siguificautlv different (at the 3-0 confidence level) from a planet model fit., We regard as acceptable any blend scenario that results in a synthetic light curve giving a $\chi^2$ for the fit that is not significantly different (at the $\sigma$ confidence level) from a planet model fit. + The brightuess variatious of the binary are eencrated using detailed calewlatious includius libi darkening.∖ gravitysepartae brighteuiug.∖⋅⊾⋅∖∖⊳⋅ reflection. ando proximity effects.," The brightness variations of the binary are generated using detailed calculations including limb darkening, gravity brightening, reflection, and proximity effects." + The properties of the candidate are tightly constrained by the spectroscopic parameters in Table 1. and were held fxed., The properties of the candidate are tightly constrained by the spectroscopic parameters in Table \ref{starproperties} and were held fixed. + The parameters of the binary compoucuts were taken from model isochrones by Carardietal.(2000).. parametrized in terius of their mass.," The parameters of the binary components were taken from model isochrones by \citet{gira2000}, parametrized in terms of their mass." + The secoudary and tertiary masses were allowed to vary over wide ranges (0.11.1 AL.) iun order to ft the Πο curve. aud the iuclination. anele was also a free parameter.," The secondary and tertiary masses were allowed to vary over wide ranges (0.1–1.4 $M_{\sun}$ ) in order to fit the light curve, and the inclination angle was also a free parameter." + By comparing the quality of the ligbit-curve fits over these ranges. we are able to constrain the properties of the secondaries. and tertiaries. that provide. acceptable fits. and we can reject other bleuds.," By comparing the quality of the light-curve fits over these ranges, we are able to constrain the properties of the secondaries and tertiaries that provide acceptable fits, and we can reject other blends." + We consider uerarchical configurations of two types: oues n which the tertiary is a star. aud ones in which it is a planet contributing no light.," We consider hierarchical configurations of two types: ones in which the tertiary is a star, and ones in which it is a planet contributing no light." + Iu the latter case he size of the planet is a free paraiucter. which we varied between 0.1 aud 2.0 Raga.," In the latter case the size of the planet is a free parameter, which we varied between 0.1 and 2.0 $R_{\rm Jup}$." + We account for he additional stars ideuti&ed by high resolution iuaging by incliding the proper amount of extra ight i our models., We account for the additional stars identified by high resolution imaging by including the proper amount of extra light in our models. + Tn all five cases we find that configurations with stellar tertiaries are mucousisteut with the light curves. for any size secoudary.," In all five cases we find that configurations with stellar tertiaries are inconsistent with the light curves, for any size secondary." + Thus. uerarchical triples iuvolving a stellar binary are ruled out.," Thus, hierarchical triples involving a stellar binary are ruled out." + When the tertiary is allowed to be a lanet. we fud that there is a range of possible solutions that vield acceptable matclics to the photometry. often tines as good as obtained roni a single star aud plauct model.," When the tertiary is allowed to be a planet, we find that there is a range of possible solutions that yield acceptable matches to the photometry, often times as good as obtained from a single star and planet model." + However. nany of those solutious cau also be excluded ou other erounds as described below.," However, many of those solutions can also be excluded on other grounds as described below." + For IKOI 152. which is a ία F dwarf. the oulv configurations consistent with the light curves involve secondary stars that are at least as lnassive as the prinaryvand therefore almost," For KOI 152, which is a mid F dwarf, the only configurations consistent with the light curves involve secondary stars that are at least as massive as the primary—and therefore almost" +"spiral galaxies such as 1100 (Rand1995;Sempere&García-Burillo1997;Calzettietal. 2005),, M551 (Calzetti 2005),, and 33147 (Casasolaetal.2008a).","spiral galaxies such as 100 \citep[][]{rand95,sempere97,calzetti05}, , 51 \citep[][]{calzetti05}, and 3147 \citep[][]{vivi08a}." +. These anti-correlations may relate to star formation efficiency and timescale variations in response to a spiral density wave., These anti-correlations may relate to star formation efficiency and timescale variations in response to a spiral density wave. +" The SF and its related tracers are often located in different regions of a galaxy: FUV emission is more prominent at the outer edge of the spiral arms, where typically dust extinction is low, while FIR emission is stronger at the inner edge."," The SF and its related tracers are often located in different regions of a galaxy: FUV emission is more prominent at the outer edge of the spiral arms, where typically dust extinction is low, while FIR emission is stronger at the inner edge." +" The FUV-CO anti-correlation found in iis particularly interesting in terms of the Kennicutt-Schmidt (KS) law (Kennicutt1998), which relates the star formation rate (SFR) density (Xsrr) to the gas surface density (Xgas)."," The FUV-CO anti-correlation found in is particularly interesting in terms of the Kennicutt-Schmidt (KS) law \citep[][]{kennicutt98}, which relates the star formation rate (SFR) density $\Sigma_{\rm SFR}$ ) to the gas surface density $\Sigma_{\rm gas}$ )." +" Bigieletal.(2008) have found that at high gas surface densities, when the gas is predominantly molecular, the KS law is linear ο with Ν 1) and a correlation between FUV and CO is expected."," \citet{bigiel08} have found that at high gas surface densities, when the gas is predominantly molecular, the KS law is linear $\Sigma_{\rm SFR}$$\sim$$\Sigma_{\rm gas}^N$, with $N$$\sim$ 1) and a correlation between FUV and CO is expected." +"λες However, more than one factor may disturb the KS correlation, especially on small spatial scales."," However, more than one factor may disturb the KS correlation, especially on small spatial scales." +" First, FUV traces older SF episodes than either oor eemission; hence the sites of potential future SF (as traced by CO) may be disconnected from past SF sites (FUV)."," First, FUV traces older SF episodes than either or emission; hence the sites of potential future SF (as traced by CO) may be disconnected from past SF sites (FUV)." +" Second, feedback from massive stellar winds and supernovae may disrupt the ISM on small spatial scales and degrade the KS correlation."," Second, feedback from massive stellar winds and supernovae may disrupt the ISM on small spatial scales and degrade the KS correlation." +" Finally, if the molecular gas is not in dynamical equilibrium, perhaps through the action of the large-scale bar, it would not be expected to be associated with sites of current SF."," Finally, if the molecular gas is not in dynamical equilibrium, perhaps through the action of the large-scale bar, it would not be expected to be associated with sites of current SF." + The lack of correlation in the context of the KS law for this and other NUGA galaxies will be discussed in a forthcoming paper dedicated to thistopic., The lack of correlation in the context of the KS law for this and other NUGA galaxies will be discussed in a forthcoming paper dedicated to thistopic. +such as that of knot C. Our method has been described in Sect.,such as that of knot C. Our method has been described in Sect. + 4.2 and. in short. is based on the computation of a large number of photoionization models with a minimal personal bias and preconceived ideas on the AGN ionizing spectrum and the gas density distribution.," \ref{details_photion} + and, in short, is based on the computation of a large number of photoionization models with a minimal personal bias and preconceived ideas on the AGN ionizing spectrum and the gas density distribution." + In particular. we also considered mixed models with combinations of density and radiation bounded clouds (B96)) as well as models with multiple density components.," In particular, we also considered mixed models with combinations of density and radiation bounded clouds \cite{B96}) ) as well as models with multiple density components." + We spanned a very wide range of model parameters which were varied randomly. and selected the relatively few (about 200) good models which came closest to reproducing the observed line ratios.," We spanned a very wide range of model parameters which were varied randomly, and selected the relatively few (about 200) good models which came closest to reproducing the observed line ratios." + The main results are summarized in Figs. 10..," The main results are summarized in Figs. \ref{figZ}," + 11 which show the distributior of element abundances required to reproduce the observed line ratios.," \ref{figZnoch} + which show the distribution of element abundances required to reproduce the observed line ratios." + The most remarkable feature is that the metal abundances are quite well constrained in spite of the very different assumptions made for the gas density distribution and shape of the AGN continua., The most remarkable feature is that the metal abundances are quite well constrained in spite of the very different assumptions made for the gas density distribution and shape of the AGN continua. + In other words. models with very different abundances fail to match the observed lines ratios in knot C regardless of the AGN spectral shape and/or gas density distribution assumed.," In other words, models with very different abundances fail to match the observed lines ratios in knot C regardless of the AGN spectral shape and/or gas density distribution assumed." + Another encouraging result is that lines from different ionization stages yield similar abundances which simply reflects the fact that the models reproduce the observed line ratios reasonably well., Another encouraging result is that lines from different ionization stages yield similar abundances which simply reflects the fact that the models reproduce the observed line ratios reasonably well. + There are however remarkable exceptions. such as the [NII[NI] and [Fell]/[FeVII] ratios which are both predicted too high.," There are however remarkable exceptions, such as the [NII]/[NI] and [FeII]/[FeVII] ratios which are both predicted too high." + Possible explanations for these differences have been discussed above (Sects. 4.5.. 4.6)).," Possible explanations for these differences have been discussed above (Sects. \ref{iron}, \ref{NII_NI}) )." + We stress here. however. that the uncertainties on |NI] have little effects on the derived nitroger abundance because N! is the most abundant ion within the partially ionized region.," We stress here, however, that the uncertainties on [NI] have little effects on the derived nitrogen abundance because $^+$ is the most abundant ion within the partially ionized region." + Therefore. the best-fit N abundance decreases by only a factor ~1.3 once the ΝΙΝΟ is increased to match the observed [NII|/[NI] ratio (cf. Fig.11)).," Therefore, the best–fit N abundance decreases by only a factor $\simeq$ 1.3 once the $^+$ $^0$ is increased to match the observed [NII]/[NI] ratio (cf. \ref{figZnoch}) )." + Note also that the He/H abundance ts only poorly constrained by the models. and although models with 70.1 are somewhat favoured. no firm conclusion about He overabundance can be drawn from the The derived abundances are summarized in Table 6 where the most striking result is the large overabundance of nitrogen relative to oxygen. +0.7 dex above the solar value. whose implications are discussed below.," Note also that the He/H abundance is only poorly constrained by the models, and although models with $>$ 0.1 are somewhat favoured, no firm conclusion about He overabundance can be drawn from the The derived abundances are summarized in Table \ref{tab_Z} + where the most striking result is the large overabundance of nitrogen relative to oxygen, +0.7 dex above the solar value, whose implications are discussed below." + An independent estimate of metallicity can be derived from the measured equivalent widths of CO stellar absorption features. using the new metallicity scale proposed and successfully applied to young LMC/SMC clusters by Oliva Origlia (1998)).," An independent estimate of metallicity can be derived from the measured equivalent widths of CO stellar absorption features, using the new metallicity scale proposed and successfully applied to young LMC/SMC clusters by Oliva Origlia \cite{oliva_origlia98}) )." + In short. the method is based on the strength of the CO(6.3) band-head at 1.62 jm whose behaviour with metallicity is modelled using synthetic spectra. of red supergiants.," In short, the method is based on the strength of the CO(6,3) band–head at 1.62 $\mu$ m whose behaviour with metallicity is modelled using synthetic spectra of red supergiants." + The equivalent width of the stellar CO lines from the central 100 x 100 pe? of Circinus are reported in Table 2 of Oliva et al. (1995)), The equivalent width of the stellar CO lines from the central 100 x 100 $^2$ of Circinus are reported in Table 2 of Oliva et al. \cite{oliva95}) ) + and yield an average metallicity of 0.73:0.3. à value remarkably close to the oxygen abundance derived above (cf.," and yield an average metallicity of $-0.7\!\pm\!0.3$, a value remarkably close to the oxygen abundance derived above (cf." + Table 6))., Table \ref{tab_Z}) ). + The nitrogen overabundance is of particular interest in view of its possible relationship with the (circumpnuclear starburst and N-enrichment from material processed through the CNO cycle., The nitrogen overabundance is of particular interest in view of its possible relationship with the (circum)nuclear starburst and N–enrichment from material processed through the CNO cycle. + According to chemical evolutionary models of starburst events. the N/O relative abundance reaches a naximum value of [N/O|~|0.6 (.e. 4 times the solar value) at about 3107 yr and remains roughly constant for several «107 years (cf.," According to chemical evolutionary models of starburst events, the N/O relative abundance reaches a maximum value of $\simeq\!+0.6$ (i.e. 4 times the solar value) at about $3\,10^8$ yr and remains roughly constant for several $\times10^8$ years (cf." + Fig., Fig. + 4 of Matteucci Padovani 1993))., 4 of Matteucci Padovani \cite{matteucci93}) ). + The nitroger overabundance mostly reflects the effect of the winds from He burning red supergiants whose surface composition is strongly N-enriched by gas dredged-up from the shell where hydrogen was burned through the CNO cycle., The nitrogen overabundance mostly reflects the effect of the winds from He burning red supergiants whose surface composition is strongly N–enriched by gas dredged–up from the shell where hydrogen was burned through the CNO cycle. + The amount and temporal evolution of the N/O abundance depends on model details. e.g. the shape of the IMF and the duration of the starburst. as well as on poorly known parameters such as the efficiency of the dredge-up and the contribution of primary N production by massive stars (e.g. Matteucci 1986)).," The amount and temporal evolution of the N/O abundance depends on model details, e.g. the shape of the IMF and the duration of the starburst, as well as on poorly known parameters such as the efficiency of the dredge–up and the contribution of primary N production by massive stars (e.g. Matteucci \cite{matteucci86}) )." + It is however encouraging to find that the observed N/O abundance (Table 6)) 1s very close to that predicted at a time which is compatible with the age of the starburst in Circinus (cf., It is however encouraging to find that the observed N/O abundance (Table \ref{tab_Z}) ) is very close to that predicted at a time which is compatible with the age of the starburst in Circinus (cf. + Fig., Fig. + 9 of Matolino et al. 1998))., 9 of Maiolino et al. \cite{maiolino}) ). + It should also be noticed that the observed absolute abundances are about an order of magnitude lower than the model predicted values. but this can be readily explained if the starburst transformed only ~LOM of the available gas into stars. in which case the chemical enrichment was diluted by a similar factor.," It should also be noticed that the observed absolute abundances are about an order of magnitude lower than the model predicted values, but this can be readily explained if the starburst transformed only $\simeq\!10\%$ of the available gas into stars, in which case the chemical enrichment was diluted by a similar factor." + This hypothesis is in good agreement with the, This hypothesis is in good agreement with the +The second parametrization has ο).(5.4.6) with s=sj» being the shortest side of the redshift-space triangle. ¢=555/54». and @ the angle between these two sides (8) and 555).,"The second parametrization has $Q_z(s,q,\theta)$ with $s=s_{12}$ being the shortest side of the redshift-space triangle, $q=s_{23}/s_{12}$, and $\theta$ the angle between these two sides $s_{12}$ and $s_{23}$ )." + Figure |. shows €Q.(s.i.0) for both our volume—limited sample (filled circles) and the Popeetal.(2004). sample (tilled stars).," Figure \ref{jingplot} shows $Q_z(s,u,v)$ for both our volume–limited sample (filled circles) and the \cite{Pope2004} sample (filled stars)." + Different panels show results for a range of triangle configurations., Different panels show results for a range of triangle configurations. + To facilitate a direct. comparison with results from theliterature. we have used the same binning scheme as Jing&Borner(1998.2004). in their analyses of the Las Campanas Redshift Survey (LCRS) and 2dF Galaxy Redshift Survey (2dFGRS).," To facilitate a direct comparison with results from theliterature, we have used the same binning scheme as \cite{JB1998,JB2004}, in their analyses of the Las Campanas Redshift Survey (LCRS) and 2dF Galaxy Redshift Survey (2dFGRS)." + The open circles show their results., The open circles show their results. + Overall. our (Qs.n.0) values are consistent with theirs. but with some obvious disagreements.," Overall, our $Q_z(s,u,v)$ values are consistent with theirs, but with some obvious disagreements." + For example. on large scales (54210b !Mpo. we find larger €.~1. while Jing&Bórner(2004. find much smaller values.," For example, on large scales $s_{12} > 10 h^{-1} {\rm Mpc}$ ), we find larger $Q_z\sim 1$, while \cite{JB2004} find much smaller values." + Although the different selection passbands of the 3dFGRS (5;) and SDSS (e. band) might account for this difference. it cannot account for the disagreement with the LCRS measurements of Jing&Bórner(1998). since the LCRS was also r—band selected.," Although the different selection passbands of the 2dFGRS $b_j$ ) and SDSS $r-$ band) might account for this difference, it cannot account for the disagreement with the LCRS measurements of \cite{JB1998} since the LCRS was also $r$ –band selected." + To quantify the disagreement. we estimated the covariances of our 3PCF estimates using the jack-knife re-sampling technique (discussed in detail in Scrantonetal.(2002). and. Zehavietal.(2002. 200555).," To quantify the disagreement, we estimated the covariances of our 3PCF estimates using the jack–knife re–sampling technique (discussed in detail in \cite{Scranton2002} and \cite{Zehavi2002,Zehavi2005}) )." +" Briefly. the jack-knife resampling technique provides an estimate of the “cosmic variance"" within a sample."," Briefly, the jack-knife resampling technique provides an estimate of the “cosmic variance” within a sample." + It is caleulated by splitting the dataset into sub-regions and then measuring the variance seen between the estimated correlation functions as sub-regions are omitted one-by-one (therefore. if there are NV subregions. there are NV’ correlation function estimates).," It is calculated by splitting the dataset into sub–regions and then measuring the variance seen between the estimated correlation functions as sub–regions are omitted one-by-one (therefore, if there are $N$ subregions, there are $N$ correlation function estimates)." +" As shown in Figure 2 of Zehavietal.(2005).. the jack-knife errors accurately reproduce the “true error"" (the dispersion measured between 100 mock galaxy catalogues). especially for the diagonal terms of the covariance matrix of the 2PCF on large scales. (for roO.5h1 Mpc. the difference between the two error estimates is always < I0)."," As shown in Figure 2 of \cite{Zehavi2005}, the jack-knife errors accurately reproduce the “true error” (the dispersion measured between 100 mock galaxy catalogues), especially for the diagonal terms of the covariance matrix of the 2PCF on large scales, (for $r>0.5h^{-1}$ Mpc, the difference between the two error estimates is always $<10\%$ )." + In what follows. we assume that the jack-knife error estimates are also accurate for the 3PCF.," In what follows, we assume that the jack-knife error estimates are also accurate for the 3PCF." + The SDSS dataset is built-up of thin “wedge-shaped” regions that are 2.5 degrees thick in declination and hundreds of degrees wide in right ascension (see York et al., The SDSS dataset is built-up of thin “wedge-shaped” regions that are 2.5 degrees thick in declination and hundreds of degrees wide in right ascension (see York et al. + 2000)., 2000). + We divided the total volume of our volume limited catalogue up into 14 sub-regions when estimating the covariance matrix., We divided the total volume of our volume limited catalogue up into 14 sub–regions when estimating the covariance matrix. + These were selected in Right Ascension along the SDSS seans., These were selected in Right Ascension along the SDSS scans. + To illustrate. Figure 2 shows one of the redshift wedges: two of the sub-regions (namely sub-regions 3 and 4) are highlighted to provide an impression of the typical size of a subregion. but also because these two particular regions will feature prominently in what follows.," To illustrate, Figure \ref{jk} shows one of the redshift wedges; two of the sub–regions (namely sub–regions 3 and 4) are highlighted to provide an impression of the typical size of a subregion, but also because these two particular regions will feature prominently in what follows." + The error bars shown in Figure |. show the diagonal elements of the covariance matrices we estimate from the jack-knife method., The error bars shown in Figure \ref{jingplot} show the diagonal elements of the covariance matrices we estimate from the jack-knife method. + The sizes of these diagonal elements (as well as the off-diagona elements) are extremely sensitive to the inclusion or exclusion of sub-regions 3 and 4., The sizes of these diagonal elements (as well as the off–diagonal elements) are extremely sensitive to the inclusion or exclusion of sub–regions 3 and 4. + This sensitivity is quantified in Figure 3 which shows the scatter between the 14 2PCFs and 3PCFs used to construct the covariance matrices., This sensitivity is quantified in Figure \ref{plotcorr} which shows the scatter between the 14 2PCFs and 3PCFs used to construct the covariance matrices. + The scatter in the 2PCFs between |2 of the l4 jack-knife datasets. which contain the supercluster seen in Figure 2.. is less than on all scales probec herein (s<40/1.+ Mpc) which is consistent with the findings of Zehavietal.(2005).," The scatter in the 2PCFs between 12 of the 14 jack–knife datasets, which contain the supercluster seen in Figure \ref{jk}, is less than on all scales probed herein $s\le +40\,h^{-1}$ Mpc) which is consistent with the findings of \cite{Zehavi2005}." +. The 2 datasets which exclude sub-regions 3 and +. have significantly different 2PCFs. up to differen on the largest scales. which is again consistent with Zehavietal.(2005) who find that this supercluster greatly affects their 2PCF on large scales and is not accounted for by their estimates of the jack—knife errors.," The 2 datasets which exclude sub–regions 3 and 4, have significantly different 2PCFs, up to different on the largest scales, which is again consistent with \cite{Zehavi2005} who find that this supercluster greatly affects their 2PCF on large scales and is not accounted for by their estimates of the jack–knife errors." +" The effect on the 3PCF of the ""Sloan Great Wall” is much greater.", The effect on the 3PCF of the “Sloan Great Wall” is much greater. + The jack-knife datasets that exclude sub-regions 3 and 4 (which contain the supercluster) differ by up to ton large scales} compared to all other 3PCFs., The jack–knife datasets that exclude sub–regions 3 and 4 (which contain the supercluster) differ by up to (on large scales) compared to all other 3PCFs. + In Figure |... we show the normalised 3PCF . for the whole dataset as well as for the datasets with sub-regions 3 and 4 excluded.," In Figure \ref{jingplot}, , we show the normalised 3PCF $Q_z$ for the whole dataset as well as for the datasets with sub–regions 3 and 4 excluded." + With the bulk of this supercluster excluded. the SDSS 3PCF has much lower €2.(s.5.0) values on large scales and is now in good agreement with the Jing&Bórner(2004). 2dFGRS 3PCF on these large scales.," With the bulk of this supercluster excluded, the SDSS 3PCF has much lower $Q_z(s,u,v)$ values on large scales and is now in good agreement with the \cite{JB2004} 2dFGRS 3PCF on these large scales." + This is also demonstrated in the error bars shown in Figure |. which were estimated using all 14 fack— datasets (dot-dashed error bars) and for the 12. jack-knifedatasets (solid error bars) which excluded the supercluster (Le.. sub-regions 3 and 4 removed).," This is also demonstrated in the error bars shown in Figure \ref{jingplot} which were estimated using all 14 jack--knife datasets (dot–dashed error bars) and for the 12 jack–knifedatasets (solid error bars) which excluded the supercluster , sub–regions 3 and 4 removed)." + As expected. the sizes of these error," As expected, the sizes of these error" +and scaled the flix so that the average of the continuum {hixes at the wavelength centers ol orders 15 and 16 agrees with our measured average value.,and scaled the flux so that the average of the continuum fluxes at the wavelength centers of orders 15 and 16 agrees with our measured average value. + We also acquired on 2004 November 23 similar NIRSPEC spectra of the A5V star 61 (νο A. These spectra were used in calibrating svnthetic models of the stellar photospheric contribution to the V836 Tau spectrum (see section 3.1)., We also acquired on 2004 November 23 similar NIRSPEC spectra of the K5V star 61 Cyg A. These spectra were used in calibrating synthetic models of the stellar photospheric contribution to the V836 Tau spectrum (see section 3.1). + The data were reduced and wavelength calibrated with the same procedure used for the V836 Tau data., The data were reduced and wavelength calibrated with the same procedure used for the V836 Tau data. + The resulting CO fIundamental spectrum is of significantly. hieher signal-to-noise than our previously reported spectrum of V836 Tau (Najila et 22003)., The resulting CO fundamental spectrum is of significantly higher signal-to-noise than our previously reported spectrum of V836 Tau (Najita et 2003). + In the earlier spectrum. CO emission was clearly detected. but the shape of the line profile was unclear and the properties of the emission were unusual.," In the earlier spectrum, CO emission was clearly detected, but the shape of the line profile was unclear and the properties of the emission were unusual." + The line widths of the low-/ Ro lines were found to be significantly narrower than the widths of the high-/ P lines (41kms! compared to 58kms 4)., The line widths of the $J$ R lines were found to be significantly narrower than the widths of the $J$ P lines $41\kms$ compared to $58\kms$ ). + In addition. the emission at 4.6//m appeared to peak at a radial velocity redward ol the stellar velocity.," In addition, the emission at $4.6\micron$ appeared to peak at a radial velocity redward of the stellar velocity." + A comparison with the new spectrum shows that this was because the blue component of the low- profiles was not recovered in the earlier data. a result of our limited ability (o correct for telluric absorption given the lower signal-to-nolse of the earlier data.," A comparison with the new spectrum shows that this was because the blue component of the $J$ profiles was not recovered in the earlier data, a result of our limited ability to correct for telluric absorption given the lower signal-to-noise of the earlier data." +" The new. higher signal-to-noise data show v=10 emission (hat is double-peaked in both the 4.6jmi and 4.9,an regions (Fig."," The new, higher signal-to-noise data show $v$ =1–0 emission that is double-peaked in both the $4.6\micron$ and $4.9\micron$ regions (Fig." + 1 and 2. respectively. top panels) and centered al the radial velocity of the star.," 1 and 2, respectively, top panels) and centered at the radial velocity of the star." + The emission equivalent width is comparable to that in the earlier spectrum., The emission equivalent width is comparable to that in the earlier spectrum. + The double-peaked line profiles of the CO v=]0 emission [rom V336 Tau are unusual in that the majority. of classical T Tauri stars Chat have been studied to date show CQO e—]0 line profiles that are centrally peakecl (Najila et 22003)., The double-peaked line profiles of the CO $v$ =1–0 emission from V836 Tau are unusual in that the majority of classical T Tauri stars that have been studied to date show CO $v$ =1–0 line profiles that are centrally peaked (Najita et 2003). + For emission arising in a disk. double-peaked profiles indicate that the emission arises [rom a limited range of disk radii.," For emission arising in a disk, double-peaked profiles indicate that the emission arises from a limited range of disk radii." + For example. this is the interpretation given to the double-peaked 2.3j/mi CO overtone (.Nv—2) emission lines that are observed from actively accreting voung stars (e.g.. WL 16. DG Tan: Carr et 11993: Najita et 11996. 2000).," For example, this is the interpretation given to the double-peaked $2.3\micron$ CO overtone $\Delta v$ =2) emission lines that are observed from actively accreting young stars (e.g., WL 16, DG Tau; Carr et 1993; Najita et 1996, 2000)." + Because the observed spectrum is a composite of the emission from (he star and clisk. some fraction of the central dip in the CO line profile may result. from absorption in the stellar photospheric component.," Because the observed spectrum is a composite of the emission from the star and disk, some fraction of the central dip in the CO line profile may result from absorption in the stellar photospheric component." + We estimated the stellar contribution to the spectrum using, We estimated the stellar contribution to the spectrum using +sample.,sample. + A main bias is due to the lack of active stars in the observed sample., A main bias is due to the lack of active stars in the observed sample. + Indeed. active stars are eliminated from samples since high activity of the star makes the detection of a plauet more dificult.," Indeed, active stars are eliminated from samples since high activity of the star makes the detection of a planet more difficult." + Ou the other hand. the effect studied here is particularly rclevaut for the most active stars with large Xrav enüssious.," On the other hand, the effect studied here is particularly relevant for the most active stars with large X–ray emissions." + We expect that CoRoT will eive us a sample of planet without biases against activity for which it will be possible to check our predictions., We expect that CoRoT will give us a sample of planet without biases against activity for which it will be possible to check our predictions. + We studied the iufiueuce of the stellar luminosity distribution ou atimospleric escape of exoplancts at orbits less than 0.1 AU., We studied the influence of the stellar luminosity distribution on atmospheric escape of exoplanets at orbits less than 0.1 AU. + We showed that a significant amount of planets cau ο evaporated over time scales of { Cyr aud that the final mass distribution is ciffereut roni the initial one., We showed that a significant amount of planets can be evaporated over time scales of 4 Gyr and that the final mass distribution is different from the initial one. + For a Jupitermass plauct with a density of 0.1 &/cur. about 5 of the planucts are lost after [Cir at an orbit of 0.02 AU.," For a Jupiter–mass planet with a density of 0.4 $^{3}$, about 5 of the planets are lost after 4 Gyr at an orbit of 0.02 AU." + For Neptunemass planets this muber increases to about 60 for a deusity of 0.8 eem., For Neptune–mass planets this number increases to about 60 for a density of 0.8 $^{3}$. + We also show that for closein orbits a large ΠΠο of planets can be croded to SuperEarth and/or may lose dense ivdrogen cuvelopes., We also show that for close–in orbits a large number of planets can be eroded to Super–Earth and/or may lose dense hydrogen envelopes. +" Further. we preseut the resulting mass distribution of au initially fat mass distribution of plaucts between 0.2 M, aud 10 Εν."," Further, we present the resulting mass distribution of an initially flat mass distribution of planets between 0.2 $_{nep}$ and 10 $_{jup}$." + For closein. orbits about 32 of the initial plaucts are lost aud the distribution is shifted to siudller Ou work is a first step to understand the evolution of plauetary mass distribution due to the stellar activity evolution., For close–in orbits about 32 of the initial planets are lost and the distribution is shifted to smaller Our work is a first step to understand the evolution of planetary mass distribution due to the stellar activity evolution. + Our approach is subject to several assunirptions that will be verified with future observatious and/or modelling., Our approach is subject to several assumptions that will be verified with future observations and/or modelling. + First of all we assume a constant density in time and a heating fuuction of 100*4., First of all we assume a constant density in time and a heating function of 100. +. Better knowledee of the deusities of exoplauets will be achieved by further modelling supported by observational evidence. while the latter assumption needs to be verified with detailed radiative trausfer calculations.," Better knowledge of the densities of exoplanets will be achieved by further modelling supported by observational evidence, while the latter assumption needs to be verified with detailed radiative transfer calculations." + The scaling law of the Ly evolution in Eq., The scaling law of the $_{X}$ evolution in Eq. + 2. is derived from a few points., \ref{scaling} is derived from a few points. + This potential source of mucertainty is however mitieated by the fact that most of the effects occur in the first Gyr (see Fie. 5)), This potential source of uncertainty is however mitigated by the fact that most of the effects occur in the first Gyr (see Fig. \ref{age}) ) + where the XN.rav evolution is better kuown., where the X–ray evolution is better known. + The UV dus aud evolution should be added in the future., The UV flux and evolution should be added in the future. + We plan to extend our work to lower mass stars where the X.rav effects may influence the habitabilitv zone., We plan to extend our work to lower mass stars where the X–ray effects may influence the habitability zone. +with respect to this parameter.,with respect to this parameter. + However. just as no dichotomy appears in the optical properties between narrow- and broad-line objects. a smooth transition between high and low X-ray luminosity ratios Is also apparent from the data.," However, just as no dichotomy appears in the optical properties between narrow- and broad-line objects, a smooth transition between high and low X-ray luminosity ratios is also apparent from the data." +" The best fit for the anti-correlation takes this form: where Hfooo is the FWHM expressed in units of 2000 km s!,"," The best fit for the anti-correlation takes this form: where $\mathrm{H\beta_{2000}}$ is the FWHM expressed in units of $2\,000$ km $^{-1}$." + We remind the reader that this correlation. which involves the soft X-ray luminosity. was calculated on the CAIXA sub-catalogue which excludes objects with a measured column density 1n excess of the Galactic one (see BOY).," We remind the reader that this correlation, which involves the soft X-ray luminosity, was calculated on the CAIXA sub-catalogue which excludes objects with a measured column density in excess of the Galactic one (see B09)." + The other correlation with H6. found in CAIXA is the one with the hard X-ray spectral slope. confirming previous claims (e.g.??)..," The other correlation with $\beta$ found in CAIXA is the one with the hard X-ray spectral slope, confirming previous claims \citep[e.g.][]{bme97,pico05}." + An anti-correlation between the 0.1-2.4 keV slope and the Hf line width was also presented by ? and later confirmed by ? and ?.., An anti-correlation between the 0.1-2.4 keV slope and the $\beta$ line width was also presented by \citet{laor94} and later confirmed by \citet{bbf96} and \citet{laor97}. + However. we note here that in CAIXA there is no significant anti-correlation between the Hf and the soft X-ray photon index.," However, we note here that in CAIXA there is no significant anti-correlation between the $\beta$ and the soft X-ray photon index." + This is likely another piece of evidence in favour of the fact that the values of E are very sensitive to the adopted models and it ts therefore better to consider the correlations between more model-independent parameters., This is likely another piece of evidence in favour of the fact that the values of $\Gamma$ are very sensitive to the adopted models and it is therefore better to consider the correlations between more model-independent parameters. + We therefore believe that our anti-correlation between the X-ray luminosity ratio and Hf. is a more basic physical correlation. because it links two independent properties of radio-quiet AGN.," We therefore believe that our anti-correlation between the X-ray luminosity ratio and $\beta$ is a more basic physical correlation, because it links two independent properties of radio-quiet AGN." + The transition from broad- to narrow-line AGN. as well as the strength of the soft excess have been often interpreted as being due to a more fundamental physical parameter. such as the accretion rate or the BH mass (e.g.?)..," The transition from broad- to narrow-line AGN, as well as the strength of the soft excess have been often interpreted as being due to a more fundamental physical parameter, such as the accretion rate or the BH mass \citep[e.g.][]{laor97}." + However. there is o significant correlation in CAIXA between Loss/Ls409 and the Eddington ratio. the BH mass or any parameter considered 1 this work other than the FWHM of Hf (a part from the one with Lj discussed above).," However, there is no significant correlation in CAIXA between $L_{0.5-2}/L_{2-10}$ and the Eddington ratio, the BH mass or any parameter considered in this work other than the FWHM of $\beta$ (a part from the one with $\Gamma_\mathrm{h}$ discussed above)." + In other words. we do not have any ndications that the strength of the soft excess may be related to any other important parameter. a part from the FWHM of Hf. ot even the accretion rate.," In other words, we do not have any indications that the strength of the soft excess may be related to any other important parameter, a part from the FWHM of $\beta$, not even the accretion rate." + The ‘classical’ interpretation of the soft excess as direct emission from the accretion dise was questioned when several studies (e.g.2?) showed that the observed temperature of resulting black body is remarkably constant across orders of magnitude of luminosities and BH masses.," The `classical' interpretation of the soft excess as direct emission from the accretion disc was questioned when several studies \citep[e.g.][]{gd04,crummy06} showed that the observed temperature of resulting black body is remarkably constant across orders of magnitude of luminosities and BH masses." + We confirmed, We confirmed +Iu the end. we retained the light curves whose characteristics are stnunarized in Table 1.,"In the end, we retained the light curves whose characteristics are summarized in Table 1." + For each target. there is a ΕΝ πο curve scement available (five segincuts in the case of the faint object SDSS J11265752). aud. iu addition. there are two NUV ποτ curve |seenmients available for SDSS J2200 0711.," For each target, there is a FUV light curve segment available (five segments in the case of the faint object SDSS J1426+5752), and, in addition, there are two NUV light curve segments available for SDSS $-$ 0741." + Figure l compares the FUV πο curves for the five tarect stars., Figure 1 compares the FUV light curves for the five target stars. + Luninositv variations are very clearly present iu the three known pulsators. SDSS 0026. SDSS J1126|5752. aud SDSS 0711. as well as in SDSS 41153|0056. a star that was not known to vary before.," Luminosity variations are very clearly present in the three known pulsators, SDSS $-$ 0026, SDSS J1426+5752, and SDSS $-$ 0741, as well as in SDSS J1153+0056, a star that was not known to vary before." + Iu contrast. photometric activity seenis to be at a iium in SDSS JOLOG|1513. another star whose variability status was unknown prior to this study.," In contrast, photometric activity seems to be at a minimum in SDSS J0106+1513, another star whose variability status was unknown prior to this study." + We find it quite remarkable. especially in view of the relative faintness of our targets. that the HST/COS conibination could pick up clear luninosity variations such as those illustrated iu Figure 1.," We find it quite remarkable, especially in view of the relative faintness of our targets, that the HST/COS combination could pick up clear luminosity variations such as those illustrated in Figure 1." + Iu order to verity the validity of our light curve extraction procedure. we first sought to recover published pulsation modes for cach of he three known variable Hot. DO stars in our sample.," In order to verify the validity of our light curve extraction procedure, we first sought to recover published pulsation modes for each of the three known variable Hot DQ stars in our sample." +" locnm Because of its Γιος», SDSS J1126]|5752 has the jt temporal coverage (15872 s) of all targets."," 1cm Because of its faintness, SDSS J1426+5752 has the best temporal coverage (15872 s) of all targets." + We show the Fourier Power Spectral Deusity of the eutire FUV data set (5 light curve segments) i the upper wart of Fieure 2., We show the Fourier Power Spectral Density of the entire FUV data set (5 light curve segments) in the upper part of Figure 2. + Clearly. the light curve is dominated * a sinele oscillation. and there is also a contribution roni its first harmonic.," Clearly, the light curve is dominated by a single oscillation, and there is also a contribution from its first harmonic." + Using standard procedures. we extracted two oscillations from the light curve. aud the characteristics of these oscillations are sununarized in the widdle of Table 2.," Using standard procedures, we extracted two oscillations from the light curve, and the characteristics of these oscillations are summarized in the middle of Table 2." + We thus find a dominant pulsation node with a period of 17.7540. s and an amplitude of 3.1640.33% of the mean brightuess10 of the star in the FUN bandpass. a result that formally corresponds to a 7.5-0 detection.," We thus find a dominant pulsation mode with a period of $\pm$ 0.40 s and an amplitude of $\pm$ of the mean brightness of the star in the FUV bandpass, a result that formally corresponds to a $\sigma$ detection." + The period corresponds remarkably well to the result of Creen et al. (, The period corresponds remarkably well to the result of Green et al. ( +2009) who found a dominant periodicity of LL7.70G6940.0008 s in the optical licht curve of SDSS J1126|5752.,2009) who found a dominant periodicity of $\pm$ 0.0008 s in the optical light curve of SDSS J1426+5752. + The FUV amplitucle of that pulsation mode is wach larger than its counterpart in the optical domain (see Table 2). but this is expected from theoretical considerations (see below) aud actually reinforces our fiudiues.," The FUV amplitude of that pulsation mode is much larger than its counterpart in the optical domain (see Table 2), but this is expected from theoretical considerations (see below) and actually reinforces our findings." + Although the second oscillation picked up in our FUV ight curve of SDSS J1126|5752 is formally ouly a 2.9-0 detection. it must be considered as “real” eiven that its veriod of 209.02+0.27 s is perfectly consistent with the value of [40.0007 s found by Caeen et al. (," Although the second oscillation picked up in our FUV light curve of SDSS J1426+5752 is formally only a $\sigma$ detection, it must be considered as “real” given that its period of $\pm$ 0.27 s is perfectly consistent with the value of $\pm$ 0.0007 s found by Green et al. (" +2009) in the optical bandpass.,2009) in the optical bandpass. + In Creen et al. (, In Green et al. ( +2009). as is he case here. that periodicity corresponds to the second üehest peak in the Fourier traustforin of the helt curve of the star. aud is associated with the first harmonic of he dominant mode.,"2009), as is the case here, that periodicity corresponds to the second highest peak in the Fourier transform of the light curve of the star, and is associated with the first harmonic of the dominant mode." + Again. the FUV amplitude of that oscillation is significantly larecr than the value of the wuplitude in the optical range.," Again, the FUV amplitude of that oscillation is significantly larger than the value of the amplitude in the optical range." + Tn the case of SDSS 0711. we combined the available FUV light curve with the two NUV liebt curves to obtain a total coverage of 7673 s. The results of our analysis are presented at the bottom of Table 2.," In the case of SDSS $-$ 0741, we combined the available FUV light curve with the two NUV light curves to obtain a total coverage of 7673 s. The results of our analysis are presented at the bottom of Table 2." +" We thus found two significant periodicities (653,170.16 s and 326.9040.0L 8) which correspond again remarkably well to the two dominaut oscillations L.397+0.056 s and 327.21840.017 3) observed in the optical light curve of that star bv Dufour et al. (", We thus found two significant periodicities $\pm$ 0.16 s and $\pm$ 0.04 s) which correspond again remarkably well to the two dominant oscillations $\pm$ 0.056 s and $\pm$ 0.017 s) observed in the optical light curve of that star by Dufour et al. ( +2009).,2009). + Specifically. in both the COS spectral range and iu the optical. the liebt curve of SDSS OFLE is dominated by an oscillation aud its first harmonic. the latter having au amplitude oulv slightly less than its parent mode.," Specifically, in both the COS spectral range and in the optical, the light curve of SDSS $-$ 0741 is dominated by an oscillation and its first harmonic, the latter having an amplitude only slightly less than its parent mode." + And the amplitudes of both periodicities are higher in the ultraviolet rauge than in the optical region., And the amplitudes of both periodicities are higher in the ultraviolet range than in the optical region. + Given that the other targets. includiue the third known pulsator SDSS 0026. have ouly a short FUV helt curve available cach for analysis. we also reported the results of another frequency extraction exercise from a," Given that the other targets, including the third known pulsator SDSS $-$ 0026, have only a short FUV light curve available each for analysis, we also reported the results of another frequency extraction exercise from a" +to the non-LTE situation by Hubeny. Hummer Lanz (1994).,"to the non-LTE situation by Hubeny, Hummer Lanz (1994)." + The calculations assumed radiative equilibrium and included the bound-free and free-free opacities of the ion and incorporated a full treatment for the blanketing effects of HI lines and the Lyman a. and > satellite opacities as computed by N. Allard (e.g. Allard et al.," The calculations assumed radiative equilibrium and included the bound-free and free-free opacities of the $^{-}$ ion and incorporated a full treatment for the blanketing effects of HI lines and the Lyman $-\alpha$ , $-\beta$ and $-\gamma$ satellite opacities as computed by N. Allard (e.g. Allard et al." + 2004)., 2004). + During the caleulation of the model structure the hydrogen line broadening was addressed in the following manner: the broadening by heavy perturbers (protons and hydrogen atoms) and electrons was treated using Allard’s data Gineluding the quasi-molecular opacity) and an approximate Stark. profile (Hubeny. Hummer Lanz 1994) respectively.," During the calculation of the model structure the hydrogen line broadening was addressed in the following manner: the broadening by heavy perturbers (protons and hydrogen atoms) and electrons was treated using Allard's data (including the quasi-molecular opacity) and an approximate Stark profile (Hubeny, Hummer Lanz 1994) respectively." +" In the spectral synthesis step detailed profiles for the Balmer lines were calculated from the Stark broadening tables of Lemke (1997),", In the spectral synthesis step detailed profiles for the Balmer lines were calculated from the Stark broadening tables of Lemke (1997). + As is our previous work (e.g. Dobbie et al., As is our previous work (e.g. Dobbie et al. + 20060). comparison between the models and the data is undertaken using the spectral fitting program XSPEC (Shafer et al.," 2006a), comparison between the models and the data is undertaken using the spectral fitting program XSPEC (Shafer et al." + 1991)., 1991). + In the present analysis all lines from to H-8 are included in the fitting process., In the present analysis all lines from $\beta$ to H-8 are included in the fitting process. + XSPEC works by folding a model through the instrument response before comparing the result to the data by means of a \> statistic., XSPEC works by folding a model through the instrument response before comparing the result to the data by means of a $\chi^{2}-$ statistic. + The best tit model representation of the data is found by incrementing free grid parameters in small steps. linearly interpolating between points in the grid. until the value of 47 is minimised.," The best fit model representation of the data is found by incrementing free grid parameters in small steps, linearly interpolating between points in the grid, until the value of $\chi^{2}$ is minimised." + Errors in the {ιδ and log g s are calculated by stepping the parameter in question away from its optimum value and redetermining minimum V until the difference between this and the true minimum 4? corresponds to lo for a given number of free model parameters (e.g. Lampton et al., Errors in the $T$$_{\rm eff}$ s and log $g$ s are calculated by stepping the parameter in question away from its optimum value and redetermining minimum $\chi^{2}$ until the difference between this and the true minimum $\chi^{2}$ corresponds to $1\sigma$ for a given number of free model parameters (e.g. Lampton et al. + 1976)., 1976). + The results of our fitting procedure are given in Table | and shown overplotted on the data in Figure I., The results of our fitting procedure are given in Table 1 and shown overplotted on the data in Figure 1. + It should be noted that the parameter errors quoted here are formal 1σ fit errors and undoubtedly underestimate the true uncertainties., It should be noted that the parameter errors quoted here are formal $1\sigma$ fit errors and undoubtedly underestimate the true uncertainties. +" In our subsequent analysis we assume more realistic levels of uncertainty of and 0.07dex in effective temperature and surface gravity respectively (e.g. Napiwotzki. Green Saffer 1999),"," In our subsequent analysis we assume more realistic levels of uncertainty of and 0.07dex in effective temperature and surface gravity respectively (e.g. Napiwotzki, Green Saffer 1999)." + V band CCD imaging of the nine white dwarfs candidate members of NGC2287 and NGC3532. was obtained on the nights of 2008/03/06 and 2008/03/07 with the Australia National University’s 40° telescope and the Wide Field Imager (WFI) located at Siding Spring Observatory., $V$ band CCD imaging of the nine white dwarfs candidate members of NGC2287 and NGC3532 was obtained on the nights of 2008/03/06 and 2008/03/07 with the Australia National University's 40” telescope and the Wide Field Imager (WFI) located at Siding Spring Observatory. + Conditions on both nights were comparatively good with photometric skies and seeing as measured from the images of ~1.5°., Conditions on both nights were comparatively good with photometric skies and seeing as measured from the images of $\sim$ 1.5”. + The WEI consists of a mosaic of 8 MIT Lincoln Labs 2048 pixel CCDs which covers an area of 52° in each pointing., The WFI consists of a mosaic of 8 MIT Lincoln Labs $\times$ 2048 pixel CCDs which covers an area of $\times$ 52' in each pointing. + Throughout our run. however. CCD6 was non-functional.," Throughout our run, however, CCD6 was non-functional." + All data were reduced using the Cambridge Astronomical Survey Unit CCD reduction toolki (Irwin Lewis 2001) to follow standard procedures. namely subtraction of the bias. flat-fielding. astrometric calibration and stacking.," All data were reduced using the Cambridge Astronomical Survey Unit CCD reduction toolkit (Irwin Lewis 2001) to follow standard procedures, namely subtraction of the bias, flat-fielding, astrometric calibration and stacking." + Aperture photometry was performed on the stacked images using a circular window with a diameter of the full width half maximum of the mean point spread function., Aperture photometry was performed on the stacked images using a circular window with a diameter of $\times$ the full width half maximum of the mean point spread function. + The Landolt fields SA98 and SAIO04 (Landolt 1992) were observec a number of times during the latter night so that instrumentamagnitudes could be transformed onto the standard V Johnson svstem. where ;1DU is a measure of the total counts from the source. {ων the exposure time and .X the airmass.," The Landolt fields SA98 and SA104 (Landolt 1992) were observed a number of times during the latter night so that instrumentalmagnitudes could be transformed onto the standard $V$ Johnson system, where $ADU$ is a measure of the total counts from the source, $t_{exp}$ the exposure time and $X$ the airmass." + The coefficients and their respective errors were determined to have the values shown in Table 2., The coefficients and their respective errors were determined to have the values shown in Table 2. +The P.V colour of each white dwarfs was estimated from the known effective temperature and surface gravity using the synthetic photometry of Bergeron et al. (, The $B-V$ colour of each white dwarfs was estimated from the known effective temperature and surface gravity using the synthetic photometry of Bergeron et al. ( +1995) as updated by Holberg Bergeron (2006).,1995) as updated by Holberg Bergeron (2006). + Our estimates of the V. magnitudes of the nine white dwarfs are listed in the tinal column of Table I., Our estimates of the $V$ magnitudes of the nine white dwarfs are listed in the final column of Table 1. + We have used the estimates of the effective temperatures. and surface gravities of the white dwarfs shown in Table | and the model grids of Bergeron et al. (, We have used the estimates of the effective temperatures and surface gravities of the white dwarfs shown in Table 1 and the model grids of Bergeron et al. ( +1995). as revised by Holberg Bergeron (2006). to derive absolute magnitudes (Mj: see Table Ij.,"1995), as revised by Holberg Bergeron (2006), to derive absolute magnitudes $_{V}$; see Table 1)." + Subsequently. we have determined the distance modulii of he nine white dwarfs. neglecting extinction. which is believed o be low AVz0.12 along the lines of sight to NGC3532 and GC2287.," Subsequently, we have determined the distance modulii of the nine white dwarfs, neglecting extinction which is believed to be low $_{V}$$\simless$ 0.12 along the lines of sight to NGC3532 and NGC2287." + These are plotted (solid bars) along with a number of distance estimates available in the literature for each of these clusters (dash-dotted lines) in Figure 2., These are plotted (solid bars) along with a number of distance estimates available in the literature for each of these clusters (dash-dotted lines) in Figure 2. + It is clear from Figure 2 hat NGC3532-6 and NGC3532-8 lie beyond NGC3532 and thus are most probably field objects., It is clear from Figure 2 that NGC3532-6 and NGC3532-8 lie beyond NGC3532 and thus are most probably field objects. + Moreover. -3 appears to lie to he foreground of NGC2287 as concluded by Koester Reimers (1981) and is also probably a field white dwarf.," Moreover, -3 appears to lie to the foreground of NGC2287 as concluded by Koester Reimers (1981) and is also probably a field white dwarf." + The remaining six objects have distance moduli which argue strongly that they are members of NGC3532 or NGC2287., The remaining six objects have distance modulii which argue strongly that they are members of NGC3532 or NGC2287. + These stars are suitable for placing constraints on the form of the IFMR., These stars are suitable for placing constraints on the form of the IFMR. + The masses and cooling times of the six cluster white dwarfs have been determined using modern evolutionary tracks supplied by the Montreal group (e.g. Fontaine. Brassard Bergeron 2001).," The masses and cooling times of the six cluster white dwarfs have been determined using modern evolutionary tracks supplied by the Montreal group (e.g. Fontaine, Brassard Bergeron 2001)." + So that this work is generally consistent with other recent studies in this area (e.g. Dobbie et al., So that this work is generally consistent with other recent studies in this area (e.g. Dobbie et al. + 20062. Williams Bolte 2007) we have adopted the ealeulations which include a mixed CO core and thick H surface layer structure.," 2006a, Williams Bolte 2007) we have adopted the calculations which include a mixed CO core and thick H surface layer structure." + The masses and cooling times shown in Table | have been derived using cubic splines to interpolate between points in this grid., The masses and cooling times shown in Table 1 have been derived using cubic splines to interpolate between points in this grid. + We note that these mass determinations are relatively insensitive to our choice of core composition and ifinstead we had adopted thin H-layer models these estimates would be systematically lower by only . . which is within our present level of precision.," We note that these mass determinations are relatively insensitive to our choice of core composition and ifinstead we had adopted thin H-layer models these estimates would be systematically lower by only $_{\odot}$ , which is within our present level of precision." +progenitor stars. which is justified because in the collapsar mocel only massive stars are LORB progenitors.,"progenitor stars, which is justified because in the collapsar model only massive stars are LGRB progenitors." + 1n our scenario HL. we assume that only massive stars (m2 mun) below some metallicity threshold. (Zc) are able to produce LOGIDs. as in the progenitor models of Hlirschietal.(2005). or Yoon.Langer&Norman(2006).," In our scenario II, we assume that only massive stars $m > m_{\rm min}$ ) below some metallicity threshold $Z_{\rm C}$ ) are able to produce LGRBs, as in the progenitor models of \citet{Hir05} or \citet*{Yoo06}." +. As explained. before. we take the mean metallicity of cold eas in cach galaxy as representative of the metallicity of the progenitor stars.," As explained before, we take the mean metallicity of cold gas in each galaxy as representative of the metallicity of the progenitor stars." + The consequences of this hypothesis will be discussed in later sections., The consequences of this hypothesis will be discussed in later sections. + In this scenario. the intrinsic LORB rate is given by Eqn.," In this scenario, the intrinsic LGRB rate is given by Eqn." +"1 for galaxies g with metallicity Z,< Zc. and is null for those with Z,2Zo.","\ref{rateint} for galaxies $g$ with metallicity $Z_g < +Z_{\rm C}$ , and is null for those with $Z_g \geq Z_{\rm C}$." + Three realizations of this scenario were computed. adopting Zo=0.1.0.3.0.6Z.. (scenarios H1. H2 and L3. respectively).," Three realizations of this scenario were computed, adopting $Z_{\rm C} = 0.1,\, +0.3,\, 0.6\, {\rm Z_{\odot}}$ (scenarios II.1, II.2 and II.3, respectively)." + Given a metallicity. threshold. we adjusted the only free parameter of our scenarios (yia) to reproduce the LORB rate measured by the BATSE experiment.," Given a metallicity threshold, we adjusted the only free parameter of our scenarios $m_{\rm min}$ ) to reproduce the LGRB rate measured by the BATSE experiment." + For this purpose. we took the olf-line GRB search of Sternetal.(2001).. which detected. 3475 GRBs with Zou2s during a [ive-time of 6.31 (70 per cent of the 9.1vr that BATSE was active). scanningvr 67 per cent of the sky.," For this purpose, we took the off-line GRB search of \citet{Ste01}, which detected 3475 GRBs with $T_{90} > 2\,{\rm s}$ during a live-time of $6.37\,{\rm yr}$ (70 per cent of the $9.1\,{\rm yr}$ that BATSE was active), scanning 67 per cent of the sky." + This translates into a LGBRD rate of Sldvr.1 above the threshold of the oll- search.," This translates into a full-sky LGRB rate of $814\,{\rm yr}^{-1}$ above the threshold of the off-line search." + The choice of this particular experiment was motivated by its good statistics and the availability of an accurate mocel for its detection efficiency., The choice of this particular experiment was motivated by its good statistics and the availability of an accurate model for its detection efficiency. + To obtain the observable rate predicted. by a given scenario. we first compute the comoving LGORB rate density [or each redshift as where V2(5005+Mpc)* is the comoving volume ol theSimulalion.," To obtain the observable rate predicted by a given scenario, we first compute the comoving LGRB rate density for each redshift as where $V = (500 \ h^{-1}\ {\rm Mpc})^3$ is the comoving volume of the." + The full-skv observed LORB rate is then where is the derivative of the comoving volume with respect to z at fixed. solid. angle. των is the maximum recshift of the simulation. d; the luminosity distance anc ff the Llubble constant at redshift 2. and paeparse the probability of detecting a LORB at redshift z with BATSE (i.e. the probability that a particular LOGRB has a peak Lux above the experiment threshold).," The full-sky observed LGRB rate is then where is the derivative of the comoving volume with respect to $z$ at fixed solid angle, $z_{\rm max}$ is the maximum redshift of the simulation, $d_L$ the luminosity distance and $H$ the Hubble constant at redshift $z$, and $p_{\rm det,BATSE}$ the probability of detecting a LGRB at redshift $z$ with BATSE (i.e. the probability that a particular LGRB has a peak flux above the experiment threshold)." + Phe integration was performed numerically. ancl the described procedure iterated over mii. until agreement with the observed rate was attained.," The integration was performed numerically, and the described procedure iterated over $m_{\rm min}$, until agreement with the observed rate was attained." + The value of pacparse(s) depends on the LORB Iuminosity Function ancl spectrum.," The value of $p_{\rm det,BATSE}(z)$ depends on the LGRB luminosity function and spectrum." + It was computed. using a Monte-Carlo scheme to simulate. at a given. redshift. a large population of LOGIUDs with different luminosities and spectra.," It was computed using a Monte-Carlo scheme to simulate, at a given redshift, a large population of LGRBs with different luminosities and spectra." + “Phen. their photon fluxes in the BATSE energy band were estimated.," Then, their photon fluxes in the BATSE energy band were estimated." + For cach LORB in this population. a second Monte-Carlo procedure rejected those events which would be undetectable. taking into account the trigger elliciencv of the Sternetal.(2001). off-line. search.," For each LGRB in this population, a second Monte-Carlo procedure rejected those events which would be undetectable, taking into account the trigger efficiency of the \citet{Ste01} off-line search." + 1 the LORB emission. were isotropic. the fraction of retained LORBs would directly. give the detection probability due to the oll-line search threshold.," If the LGRB emission were isotropic, the fraction of retained LGRBs would directly give the detection probability due to the off-line search threshold." + To account for the beamed emission of LORBs (e.g.οποίοιet.al.2005)... and assuming that the distribution of jet opening angles is independent of the LORB luminosity. spectrum and redshift. we obtain paapvrsk by multiplying this fraction by the mean beaming fraction of the jets pi=107.," To account for the beamed emission of LGRBs \citep[e.g.][]{Yon05}, and assuming that the distribution of jet opening angles is independent of the LGRB luminosity, spectrum and redshift, we obtain $p_{\rm det,BATSE}$ by multiplying this fraction by the mean beaming fraction of the jets $p_{\rm jet} = 10^{-3}$." + The luminosity function ancl spectral parameters distributions were taken from Daigneetal. (2006).., The luminosity function and spectral parameters distributions were taken from \citet{Dai06}. . + This procedure. not. only allows us to mimic the observational process but also has. provided. us with estimates of the minimum mass for the progenitor stars., This procedure not only allows us to mimic the observational process but also has provided us with estimates of the minimum mass for the progenitor stars. +" The values of mai, obtained. in solar masses. are 91.4d:0.1. 3:13.00.2. 44.3404. anc 76.020.3 for scenarios L. 1.1. H2 ancl 1.3. respectively: the uncertainties rellect. Poissonian errors in the number of observed LORBs (see also Table 1)."," The values of $m_{\rm min}$ obtained, in solar masses, are $91.4 \pm 0.1$, $13.9 \pm 0.2$, $44.3 \pm 0.4$, and $76.0 \pm 0.3$ for scenarios I, II.1, II.2 and II.3, respectively; the uncertainties reflect Poissonian errors in the number of observed LGRBs (see also Table 1)." + A meaningful definition of a LORB host galaxy in our scenarios is not as straightforward as it might. seem at first sight., A meaningful definition of a LGRB host galaxy in our scenarios is not as straightforward as it might seem at first sight. + The naivve definition of a host galaxy. at a given redshift as being any galaxy g with a LOBR. rate of ronalg.2)z0 (or identically. SER(g.z)- 0) is not useful because the resulting host population would include galaxies with arbitrarily low LORB rates.," The naïvve definition of a host galaxy, at a given redshift $z$, as being any galaxy $g$ with a LGBR rate of $r_{\rm GRB}(g,z) > 0$ (or identically, ${\rm SFR}(g,z) > 0$ ) is not useful because the resulting host population would include galaxies with arbitrarily low LGRB rates." + Low rates are. better understood. in statistical terms. as very low probabilities of producing a LORB per unit time.," Low rates are better understood in statistical terms, as very low probabilities of producing a LGRB per unit time." + Fhis implies that the corresponding galaxies have a low probability. of being observed. as host galaxies., This implies that the corresponding galaxies have a low probability of being observed as host galaxies. + Hence. this naivve definition would generate a host galaxies sample biased to low SER galaxies.," Hence, this naïvve definition would generate a host galaxies sample biased to low SFR galaxies." + To fix this problem. the definition of a host ealaxy could. be based on the number of LORBs IN(g.2)=rong.>)A? produced in cach galaxy during a time interval AM in its rest frame. so that host galaxies are only those with N(g.c)c1. as in Campisietal.(2009).," To fix this problem, the definition of a host galaxy could be based on the number of LGRBs $N(g,z) = r_{\rm GRB}(g,z) \Delta +t$ produced in each galaxy during a time interval $\Delta t$ in its rest frame, so that host galaxies are only those with $N(g,z) \geq 1$, as in \citet{Cam09}." +.. Llowever this cut-olf. and the resulting population. would be dependent on the more or less arbitrary choice of Af. Given the above arguments. we preferred. instead. a probabilistic approach. defining the likelihood of a galaxy being observed as aLGICUD host.," However this cut-off, and the resulting population, would be dependent on the more or less arbitrary choice of $\Delta t$ Given the above arguments, we preferred instead a probabilistic approach, defining the likelihood of a galaxy being observed as aLGRB host." + This likelihood is then used as à weight to compute the properties of the observable host ealaxies sample. which would be in this way comparable to," This likelihood is then used as a weight to compute the properties of the observable host galaxies sample, which would be in this way comparable to" +For each of the three selected regions (200x300 pixels area) as marked on Figures 1.. the light curves for average intensities were made.,"For each of the three selected regions $\times$ 300 pixels area) as marked on Figures \ref{fig1}, the light curves for average intensities were made." + This gives. à set of 10000 light curves while averaging on the intensities of the smaller 3x pixels regions. 3600 light curves while averaging on the intensities of smaller 5x pixels regions. and 1089 light curves while the selected smaller regions are found to be 9x pixels.," This gives, a set of 10000 light curves while averaging on the intensities of the smaller $3\times3$ pixels regions, 3600 light curves while averaging on the intensities of smaller $5\times5$ pixels regions, and 1089 light curves while the selected smaller regions are found to be $9\times9$ pixels." + These light curves were fed to the network., These light curves were fed to the network. + The network labeled each light curve with a set of individual kev parameters, The network labeled each light curve with a set of individual key parameters +Because seandiuii uuderabundance. alone with the oue of caleuu. characterizes AmFui stars. i is a kev clement to better understand the AmFin pWLOMCHOM.,"Because scandium underabundance, along with the one of calcium, characterizes AmFm stars, it is a key element to better understand the AmFm phenomenon." + Scandi Is du alk lonisatkn state. with noble gas configuration in lavers clo«x| to the bottoli of the II convection Zone. and close to the botOni of the Fe convection ZOje predicteL by. Richer. Aichaud Turctte (2000). mut not at the same posikns as for the correspoudine caleiuni colifieurations.," Scandium is in an ionisation state with noble gas configuration in layers close to the bottom of the H convection zone, and close to the bottom of the Fe convection zone predicted by Richer, Michaud Turcotte (2000), but not at the same positions as for the corresponding calcium configurations." + Aiv model of AimFin stars shoud be able te» explain uuderabuualicos for both elements., Any model of AmFm stars should be able to explain underabundances for both elements. + Therefore. the studv of Se diSIO could possibly sliec| light on which one of the two diUNIO scenarios. cüffusioi below the II convection. zoue while juchiding amass loss (Alecian 1996: rereatter Ay). aud diffusion below the Fe convection zoje (Richer. Alichaud Turcotte 2000: hereafter BY). is more likely to explain the abuudauce anolalies of tjose stars.," Therefore, the study of Sc diffusion could possibly shed light on which one of the two diffusion scenarios, diffusion below the H convection zone while including mass loss (Alecian 1996; hereafter ), and diffusion below the Fe convection zone (Richer, Michaud Turcotte 2000; hereafter ), is more likely to explain the abundance anomalies of these stars." + Iu Sec. 2.," In Sec. \ref{sec:grad}," + we preseut Grad Calculations for Se while using au interpolation 1nethod related o the new parametric foruulae described in Pavers I aud IL, we present $g_{rad}$ calculations for Sc while using an interpolation method related to the new parametric formulae described in Papers I and II. + In Sec. H3..," In Sec. \ref{sec:abundances}," +" we discuss our Grad With respect to the observed abundance anemales of Ca aud Sc in AmFu stars,", we discuss our $g_{rad}$ with respect to the observed abundance anomalies of Ca and Sc in AmFm stars. + A short conclusion follows., A short conclusion follows. + Receutly. aniuproved parametric nmiethod for calculating Grad at large «yptical depths weis developed. (Papers I aud ID.," Recently, an improved parametric method for calculating $g_{rad}$ at large optical depths was developed (Papers I and II)." + The muprovenieuts. as conpared to the method of Alecian Ariu (19908.b). were made possible because of the use of OP opaciies and data (Seaton et al.," The improvements, as compared to the method of Alecian Artru (1990a,b), were made possible because of the use of OP opacities and data (Seaton et al." + 1992)., 1992). + Ouce the parameters of the various lous are calculated. the parametric method las the advantage of ceiving relatively accurate gig without having o deal with the enonuous aluotut of atomic data normalvo necessary to obtain grad," Once the parameters of the various ions are calculated, the parametric method has the advantage of giving relatively accurate $g_{rad}$ without having to deal with the enormous amount of atomic data normally necessary to obtain $g_{rad}$." + The wraleters for 12 trace eclients (C. δν O. Ne. Na. Me. Al. Si. S. Ay. Ca axd Fe) were calculated in Paper II.," The parameters for 12 trace elements (C, N, O, Ne, Na, Mg, Al, Si, S, Ar, Ca and Fe) were calculated in Paper II." + The method usec here to calculate the gag of Sc is base onu this parametric foru of grag formulae (SVP method) of both botnud-bouuc (Alecian 1985. Alecian Aytii 1990a. Papers I aud II) aud. bound-free CAleciau 1991. Papers I an 11) traαμΊο».," The method used here to calculate the $g_{rad}$ of Sc is based on this parametric form of $g_{rad}$ formulae (SVP method) of both bound-bound (Alecian 1985, Alecian Artru 1990a, Papers I and II) and bound-free (Alecian 1994, Papers I and II) transitions." + lu this inethod. radiaive accelerations due to atouic lines depend on 1 parameters for cach ion (Eq. {," In this method, radiative accelerations due to atomic lines depend on 4 parameters for each ion (Eq. (" +1) of Paper ID). calculated at lavers where the relative opulatiou of the ion / is near its maxima.,"1) of Paper II), calculated at layers where the relative population of the ion $i$ is near its maximum." + The first ]oxwuueter o; (Eqs. (, The first parameter $\phi_{i}^{*}$ (Eqs. ( +10) to (13) of Paper D) is related to he strength of the bound-bound transitions through a weighed average of the gf values.,10) to (13) of Paper I) is related to the strength of the bound-bound transitions through a weighted average of the $gf$ values. + The second parameter c6;; (see Eq. (, The second parameter $\psi_{i}^{*}$ (see Eq. ( +10) of Paper Ij) is related to the average width of the line profiles aud controls saturation.,14) of Paper I) is related to the average width of the line profiles and controls saturation. + A third piriuueer & is related to the ion contribution to the total opacivy (we will neglect it for Sc since it is of little importance for less abuidaut elements)., A third parameter $\xi_{i}^{*}$ is related to the ion contribution to the total opacity (we will neglect it for Sc since it is of little importance for less abundant elements). + Meauwhile. the value of the parameter a; is determine by fitting our parauetric accelerations to hose obtained by a 1iore accurate moethod (Seaton 1997. see Paper I for more details)," Meanwhile, the value of the parameter $\alpha_i$ is determined by fitting our parametric accelerations to those obtained by a more accurate method (Seaton 1997, see Paper I for more details)." + A sinular xuranetrie equation was also used for bound-yee transitiois (Eq. (, A similar parametric equation was also used for bound-free transitions (Eq. ( +9) of Paper ID.,9) of Paper II). + This approximate ornmula is less accurate than for bouud-bound trausitious and needs the kuowledee of enerev levels for cach ion aud the courtation of the partition functions at cach uodel laver., This approximate formula is less accurate than for bound-bound transitions and needs the knowledge of energy levels for each ion and the computation of the partition functions at each model layer. + Two other parameters a; aud b; (correction xuueters for bouud-free acceleration) axe determined for each ion by a same kind of fittii8 procedure as employed or lines., Two other parameters $a_{i}$ and $b_{i}$ (correction parameters for bound-free acceleration) are determined for each ion by a same kind of fitting procedure as employed for lines. + These six parameters. along with the formulae of Paper IT. define the SVP imehod.," These six parameters, along with the formulae of Paper II, define the SVP method." +" They are caleulated rere, dn au Au stelay model. or the 12 trace elements ueutioned above."," They are calculated here, in an Am stellar model, for the 12 trace elements mentioned above." + These eleucnuts are treated by Seatou (1997). aud our fitting procelure is achieved using version OPCD 2.1 of the tables avaialeat CDS?.," These elements are treated by Seaton (1997), and our fitting procedure is achieved using version OPCD 2.1 of the tables available at." +. The accurate evaluatioll of radiative accelerations liocessiates detailed aud coulete atomic data Guchiding all stroic transitions) for eacl clement., The accurate evaluation of radiative accelerations necessitates detailed and complete atomic data (including all strong transitions) for each element. + This is still lacking for scaidimm., This is still lacking for scandium. + A rather Lue5 data base for enerev levels can be iud iu the NIST Ate)nic SpectraDatabase. but this daabase is incomplete tx bouud-bound trausitions. especiaIv iu veeards to transiion probabilities.," A rather large data base for energy levels can be found in the NIST Atomic Spectra, but this database is incomplete for bound-bound transitions, especially in regards to transition probabilities." + As shown in Alecian Artru (1990b). piarineters o; and c; generally vary quite monotonously along isoclectronic sequences.," As shown in Alecian Artru (1990b), parameters $\phi_{i}^{*}$ and $\psi_{i}^{*}$ generally vary quite monotonously along isoelectronic sequences." + It is tjen possible to approximate these parameters for some of the iois of the eleucuts for which atomic data are lacking. bv iuterpoating between the nearest isoclectronic ucigibours with kuownu pirneters. to the ious 1i question.," It is then possible to approximate these parameters for some of the ions of the elements for which atomic data are lacking, by interpolating between the nearest isoelectronic neighbours with known parameters, to the ions in question." + The interpolation is done via the atomic charge variale., The interpolation is done via the atomic charge variable. +" We have then considered that to apply tus interpolation method to Sc 10115, usine well calculated parameters from isoclecronic ucighbours. is preferable han to calculate civectly the parameters. using incomplete atomic data."," We have then considered that to apply this interpolation method to Sc ions, using well calculated parameters from isoelectronic neighbours, is preferable than to calculate directly the parameters, using incomplete atomic data." + Also. it would be mupossible to do the fitting procedure o| Paper I. since the g4;4 of Sc are not included in Seato1 (1997) or elsewhere.," Also, it would be impossible to do the fitting procedure of Paper II, since the $g_{rad}$ of Sc are not included in Seaton (1997) or elsewhere." + We have interpolated the parameters for he bouixd-bound transitions of the ious ScIII to ScXX. mt not for ScII since paraiueters for a1 isoelectronic neighbour are nissug (nuaucly the parameters for Cal).," We have interpolated the parameters for the bound-bound transitions of the ions ScIII to ScXX, but not for ScII since parameters for an isoelectronic neighbour are missing (namely the parameters for CaI)." + The xwanmieters which were directly evaluated by Aleciau Artru (10905) for ScIT youn atomic data were then used., The parameters which were directly evaluated by Alecian Artru (1990b) for ScII from atomic data were then used. + Parameters e; aud b; were respectively se to laud 0 (this corres»»uds to no correction)., Parameters $a_{i}$ and $b_{i}$ were respectively set to 1 and 0 (this corresponds to no correction). + Relative populations of Se jous and partition fuwctions have been calculated usine the atomic cnerev level data of the NIST Atomic Spectra Daalbasce., Relative populations of Sc ions and partition functions have been calculated using the atomic energy level data of the NIST Atomic Spectra Database. +On sub-areseconcl scales. the central radio components. or ‘cores’. of radio galaxies are understood. as a result of ΝΕΟΙ observations. to be the unresolved. bases of the relativistic jets that. transport. energy. to the lobes. seen. in partly self-absorbed: svnchrotron radiation.,"On sub-arcsecond scales, the central radio components, or `cores', of radio galaxies are understood, as a result of VLBI observations, to be the unresolved bases of the relativistic jets that transport energy to the lobes, seen in partly self-absorbed synchrotron radiation." + Superluminal motion is seen in some sources. and the supposed unification of racio galaxies with BL Lac objects and quasars UlUrry Padovani 1995) requires the cores to be relativisticallv beamecl with Lorentz factors zz:5.," Superluminal motion is seen in some sources, and the supposed unification of radio galaxies with BL Lac objects and quasars Urry Padovani 1995) requires the cores to be relativistically beamed with Lorentz factors $\ga 5$." + Larcleastle Worrall (1999) and Canosa ((1999) have found. that the nuclear soft-N-ray emission in radio galaxies is well correlated. with the core radio emission., Hardcastle Worrall (1999) and Canosa (1999) have found that the nuclear soft-X-ray emission in radio galaxies is well correlated with the core radio emission. + This finding implies that the N-ray emission is also relativistically beamec ancl so must originate in the jet. ancl is qualitatively consistent with models in which the strong X-ray emission of BL Lac objects is largely a resul of relativistic boosting.," This finding implies that the X-ray emission is also relativistically beamed and so must originate in the jet, and is qualitatively consistent with models in which the strong X-ray emission of BL Lac objects is largely a result of relativistic boosting." + A snapshot survey of 3CR. radio sources with theL157 Wide Field and Planetary Camera 2 (WEDPC2) (De oll 11996) has detected. unresolved. optical nuclear components in a large number of objects. particularly at low redshift (Alartel 11999). where. because of the Iuminositv-redshift correlation in the 3CTU sample. most radio sources are of tvpe EIU rather than PRI (Fanarolf Riley 1974).," A snapshot survey of 3CR radio sources with the Wide Field and Planetary Camera 2 (WFPC2) (De Koff 1996) has detected unresolved optical nuclear components in a large number of objects, particularly at low redshift (Martel 1999), where, because of the luminosity-redshift correlation in the 3CR sample, most radio sources are of type FRI rather than FRII (Fanaroff Riley 1974)." + Chiaberge. Capetti Celotti (1999) find that the optical luminosities of the FRI radio galaxies are correlated: with the Iuminosities of their 5-Gllz radio cores. which. by the same argument as used by Hardceastle Worrall for the X-ray emission. implies a jet-relatec origin for the nuclear optical emission.," Chiaberge, Capetti Celotti (1999) find that the optical luminosities of the FRI radio galaxies are correlated with the luminosities of their 5-GHz radio cores, which, by the same argument as used by Hardcastle Worrall for the X-ray emission, implies a jet-related origin for the nuclear optical emission." + Capetti Celotti (1999) compare the optical nuclear Iuminosities of five of the detected PRI radio galaxies with those of BL Lac objects matched in isotropic properties and find support or unification in the optical assuming relativistic jets with Lorentz [actors in the range 5 to 10., Capetti Celotti (1999) compare the optical nuclear luminosities of five of the detected FRI radio galaxies with those of BL Lac objects matched in isotropic properties and find support for unification in the optical assuming relativistic jets with Lorentz factors in the range 5 to 10. + In this paper we examine he radio. optical and X-ray nuclear Πάνος and bIuminosities of a sample of Iow-redshift radio galaxies in order to explore he X-ray and optical emission mechanisms.," In this paper we examine the radio, optical and X-ray nuclear fluxes and luminosities of a sample of low-redshift radio galaxies in order to explore the X-ray and optical emission mechanisms." + Throughout the paper we adopt a cosmology with kms ! Mpe+., Throughout the paper we adopt a cosmology with $H_0 = 50$ km $^{-1}$ $^{-1}$. + Spectral index a is defined in the sense hat [lux density is proportional to £mm, Spectral index $\alpha$ is defined in the sense that flux density is proportional to $\nu^{-\alpha}$. + We selected the 27 radio galaxies with z«0.06 from the 3CRI sample (Laing. Riley Longair 1983. Laing Riley. in prep.):," We selected the 27 radio galaxies with $z<0.06$ from the 3CRR sample (Laing, Riley Longair 1983, Laing Riley, in prep.);" + the redshift limit was chosen to ensure &ood spatial resolution. so that nuclei could adequately be separated from the clusty discs that surround. them.," the redshift limit was chosen to ensure good spatial resolution, so that nuclei could adequately be separated from the dusty discs that surround them." + The status of the archival WEDPC?2257 observations of these objects is tabulated in Table. 1.., The status of the archival WFPC2 observations of these objects is tabulated in Table \ref{ss}. + Phe majority of the observations were made as part of the snapshot. survey. ancl so are short (280 s) ancl use the Planetary Camera (PC) with the E702W filter. but we have used longer observations in," The majority of the observations were made as part of the snapshot survey, and so are short (280 s) and use the Planetary Camera (PC) with the F702W filter, but we have used longer observations in" +quasars. and Fourier power spectrum analysis for a 15 vear homogeneous subset ofthe data confirms that there is more power on all timescales in the blue passband: compared with red.,"quasars, and Fourier power spectrum analysis for a 15 year homogeneous subset of the data confirms that there is more power on all timescales in the blue passband compared with red." + This ellect. may at least in part. be due to the contribution of red light from the underlying host. galaxy. and can be corrected. for in such a way that the variation comes close to being achromatic for all quasars.," This effect may at least in part be due to the contribution of red light from the underlying host galaxy, and can be corrected for in such a way that the variation comes close to being achromatic for all quasars." + Current theories of accretion disc instability can explain colour changes seen in observations of Sevíert galaxies., Current theories of accretion disc instability can explain colour changes seen in observations of Seyfert galaxies. + The standard optically thick accretion disk is characterised. by a radial temperature gradient which when perturbed will tend to exhibit changes in integrated colour., The standard optically thick accretion disk is characterised by a radial temperature gradient which when perturbed will tend to exhibit changes in integrated colour. + However. they do not appear to be able to produce the smearing out in the red of blue features which is a characteristic of many quasar light curves.," However, they do not appear to be able to produce the smearing out in the red of blue features which is a characteristic of many quasar light curves." + An alternative explanation for the observed. variability of quasars is that it is caused. by. the microlensing cllects of a population of planctary or sub-stellar mass bodies distributed. along the line of sight., An alternative explanation for the observed variability of quasars is that it is caused by the microlensing effects of a population of planetary or sub-stellar mass bodies distributed along the line of sight. + In. this case achromatic variation can be accounted for by microlensing of an optically thin (advection dominated) disc., In this case achromatic variation can be accounted for by microlensing of an optically thin (advection dominated) disc. + Chromatic variations are possible with a conventional optically thick disc with a radial colour graclicnt. which is large compared with the Einstein radius of the lenses.," Chromatic variations are possible with a conventional optically thick disc with a radial colour gradient, which is large compared with the Einstein radius of the lenses." + This can also explain the the observed smearing out in the red of eusp-like features in the blue., This can also explain the the observed smearing out in the red of cusp-like features in the blue. + ] wish to thank Alan Lleavens for suggestions with regard to the subtraction of the [ux of the underlying host galaxy., I wish to thank Alan Heavens for suggestions with regard to the subtraction of the flux of the underlying host galaxy. +than approximately |LO2108.6.,than approximately $+40\;21\;08.6$. +theThis difference may be due to the incorrect alignment of N-ray and radio data. the poor S/N of the lower contour levels. the cussion from the individual stellar winds extending further from the stars than expected. or because our models do not adequately represent the actual ciission from the wind collision region.," This difference may be due to the incorrect alignment of the X-ray and radio data, the poor S/N of the lower contour levels, the emission from the individual stellar winds extending further from the stars than expected, or because our models do not adequately represent the actual emission from the wind collision region." + The κα ο maps preseuted iu this paper show evidence for thefirst direct spatially resolved) detection of N-rav. cussion from the wind-wind collision gone of a dassive early-type binary., The HRC-I X-ray maps presented in this paper show evidence for the direct spatially resolved) detection of X-ray emission from the wind-wind collision zone of a massive early-type binary. + The data from WR 117 has a EWIIM of ©0.8”. roughly double that of the PSF.," The data from WR 147 has a FWHM of $\approx 0.8\arcsec$, roughly double that of the PSF." + Models of the wind collision zoue are also in rough quantitative agreement. aud im particular predict a FWIIM consisteut with that measured.," Models of the wind collision zone are also in rough quantitative agreement, and in particular predict a FWHM consistent with that measured." + To our knowledec this is the first time that N-ray enissiou has heen resolved ina stellar svsteni., To our knowledge this is the first time that X-ray emission has been resolved in a stellar system. + A future exposure of longer duration is necessary to accurately align the Nay. aud. optical/radio frames. to potentially constrain the inclination and wind mniomeutuui ratio. to determune if the angular scale of the observed Cluission is consistent with models of the wind collision. aud to determine the exact coutribution of the individual stars to the ταν CUSSION.," A future exposure of longer duration is necessary to accurately align the X-ray and optical/radio frames, to potentially constrain the inclination and wind momentum ratio, to determine if the angular scale of the observed emission is consistent with models of the wind collision, and to determine the exact contribution of the individual stars to the X-ray emission." +" Siuce the shocked WN& wind dominates the colliding winds X-ray emission. it is uulikelv that separate values for Mop aud cx,,,, can be determined from X-ray. observations of this svstein."," Since the shocked WN8 wind dominates the colliding winds X-ray emission, it is unlikely that separate values for $\Mdot_{\rm OB}$ and $v_{\infty_{\rm OB}}$ can be determined from X-ray observations of this system." + Finally. we uote the discovery that there is sometimes a discrepancy between estimates of the X-ray Iuuinosity ratio οἳ the shocked winds computed using the equations in Usov (1992)) aud calculations directly from wdrodvuamuc simulations.," Finally, we note the discovery that there is sometimes a discrepancy between estimates of the X-ray luminosity ratio of the shocked winds computed using the equations in Usov \cite{U1992}) ) and calculations directly from hydrodynamic simulations." + While the ease of application of the elegaut equatious in Usov 19923) have led to their widespread use. thisissue is examined iu iore detail in a urther paper (Pittard Steveus 2002)].," While the ease of application of the elegant equations in Usov \cite{U1992}) ) have led to their widespread use, this issue is examined in more detail in a further paper (Pittard Stevens \cite{PS2002}) )." +Several high time resolution instruments are able to measure the arrival time of visible photons with an internal precision in the range LO - 100 picoseconcds. and refer those arrivals to the commonly used. UTC scale with an accuracy. of the order of 500. picoseconds (see for instance. Darbieri ct al.,"Several high time resolution instruments are able to measure the arrival time of visible photons with an internal precision in the range 10 - 100 picoseconds, and refer those arrivals to the commonly used UTC scale with an accuracy of the order of 500 picoseconds (see for instance, Barbieri et al." + 2010 and Naletto Barbieri. 2009)). The step of referring the arrival times to UTC is usually done with the intermediary of time signals broacdcasted over radio [frequencies (e.g. GPS. GLONASS. Galileo GNSS). signals which are very accurately corrected. [for atmospheric propagation effects like ionospheric scintillation and wet troposphere refraction.," \cite{barbieri} and Naletto Barbieri, \cite{naletto}) The step of referring the arrival times to UTC is usually done with the intermediary of time signals broadcasted over radio frequencies (e.g. GPS, GLONASS, Galileo GNSS), signals which are very accurately corrected for atmospheric propagation effects like ionospheric scintillation and wet troposphere refraction." + Lo the usual astronomical applications. the arrival times of optical photons are not corrected to the same The possibility to perform such correction is actually shown bv accurate laser ranging to geodetic satellites (recall that 1 nanosecond in vacuum corresponds to 30 em)," In the usual astronomical applications, the arrival times of optical photons are not corrected to the same The possibility to perform such correction is actually shown by accurate laser ranging to geodetic satellites (recall that 1 nanosecond in vacuum corresponds to 30 cm)." + For instance. Ixral et al. (2005))," For instance, Kral et al. \cite{ivan}) )" + quote a precision. of [ew picoseconds by taking into account the atmospheric seeing., quote a precision of few picoseconds by taking into account the atmospheric seeing. + Aloreover a recent. discussion has been performed. by Duchy D.Wijava and Fritz Ix. Brunner (20112) on the atmospheric range correction for two-Lrequency Satellite Laser Ranging (SLI) Alotivated by our own very. precise time measurements on celestial sources with Aqueve ancl Iqueye. (9010 and Naletto Darbieri.. 2009)). we have undertaken the calculation of the delay. and. delay dispersion of visible photon arrival times in the usual conditions prevailing in astronomical observatories.," Moreover a recent discussion has been performed by Dudy D.Wijaya and Fritz K. Brunner \cite{wijaya}) ) on the atmospheric range correction for two-frequency Satellite Laser Ranging (SLR) Motivated by our own very precise time measurements on celestial sources with Aqueye and Iqueye \cite{barbieri} and Naletto Barbieri, \cite{naletto}) ), we have undertaken the calculation of the delay and delay dispersion of visible photon arrival times in the usual conditions prevailing in astronomical observatories." + In the first. step of our procedure. the Marini-Murrayv. model (Marini Murray. 1973)) is used to caleulate the correction AR to the/ optical path of photons in Through this caleulation. the atmospheric refractive index » and a fixed delay time independent of the photometric night quality is derived.," In the first step of our procedure, the Marini-Murray model (Marini Murray, \cite{marini}) ) is used to calculate the correction $\Delta R$ to the optical path of photons in Through this calculation, the atmospheric refractive index $n$ and a fixed delay time independent of the photometric night quality is derived." + Phen the photon paths are correlated with the astronomical Vhrough this relation we cerive a statistical κο of delay times as function of the Fried radius ro and telescope diameter., Then the photon paths are correlated with the astronomical Through this relation we derive a statistical set of delay times as function of the Fried radius $r_{0}$ and telescope diameter. + Finally. the dillerence of these delay times with the fixed delay gives the Reversing this model we also developed a theoretical mathematical model for. the ry calculating through the observation of these The results are expounded in the present paper. which is organized. as follows:," Finally, the difference of these delay times with the fixed delay gives the Reversing this model we also developed a theoretical mathematical model for the $r_{0}$ calculating through the observation of these The results are expounded in the present paper, which is organized as follows:" +"ο eas coniponeuts correspondiug to REC aud BEC of IL,.",circumstellar gas components corresponding to REC and BEC of $_\alpha$. + The effec of gaseous accretion are probably also seen on the nebular eimissiou lines., The effect of gaseous accretion are probably also seen on the nebular emission lines. +" We fouud a correlation between the total equivalent width in IL, and the peak uchulay emission iu IL, aud Na L Figures d 2 show that the nebular IL, and Na I lines were of τμήμα! iuteusitv. (1.06 above the contiuuuau) ou 7th Oct. The same lines are strongest (IL, ~7.9. on the blue wine of the REC. and Na I D2 line ~3.0) ou 27th Oct. when the IL, equivalent width is masini (-LO0A))."," We found a correlation between the total equivalent width in $_\alpha$ and the peak nebular emission in $_\alpha$ and Na I. Figures 1 2 show that the nebular $_\alpha$ and Na I lines were of minimal intensity $\sim$ 1.06 above the continuum) on 7th Oct. The same lines are strongest $_\alpha$ $\sim$ 7.9, on the blue wing of the REC, and Na I D2 line $\sim$ 3.0) on 27th Oct. when the $_\alpha$ equivalent width is maximum )." + We plan to model the impact of the episodic accretion ou the nebula using a ploto-ionizing code (for ec., We plan to model the impact of the episodic accretion on the nebula using a photo-ionizing code (for eg. +" the Cloudy code) which will be published along with the cutive spectra of LIT, 231 clsewhere in the near future.", the Cloudy code) which will be published along with the entire spectra of $_\alpha$ 234 elsewhere in the near future. + Figure 2 shows the variability in the Na I D2 and D1 lines., Figure 2 shows the variability in the Na I D2 and D1 lines. +" Redshifted Absorption Compoucuts (RAC) were observed ta Na I lines ou only one uieht (13 Oct. 2003) with a maximaun redshitted velocity of 200 κής, The RACs associated with both the Na I lines are of simular depth (for redshüfts << LOOkmi/s) which is au iudicator of saturation (unshielded Na I cola deusity = Lol fem? (de Wiuter et al.", Redshifted Absorption Components (RAC) were observed in Na I lines on only one night (13 Oct. 2003) with a maximum redshifted velocity of 200 km/s. The RAC's associated with both the Na I lines are of similar depth (for redshifts $\leq$ 100km/s) which is an indicator of saturation (unshielded Na I column density = $^{12}$ $^2$ )(de Winter et al. + 1999)., 1999). + However. the imiaxinua depth ofthese compoucuts is about of the stellar coutiuuun miplviug partial coverage of the stellar photosphere.," However, the maximum depth of these components is about of the stellar continuum implying partial coverage of the stellar photosphere." + Iu addition to Na I RAC. very uuld BEC (60 to 100 kiu/s aud. 1.05 to 1.07 above the continuum) are also seen on the 13th Oct. spectra.," In addition to Na I RAC, very mild BEC (60 to 100 km/s and 1.05 to 1.07 above the continuum) are also seen on the 13th Oct. spectra." +" The variations in Ie I and IL, are observed even when no Na I RAC BEC are seen (see figures 1 2).", The variations in He I and $_\alpha$ are observed even when no Na I RAC BEC are seen (see figures 1 2). +" Thus the appearance and disappearance of the Na TRAC BEC ave uucorrelated with the variations in IL, aud Ile I liue profiles.", Thus the appearance and disappearance of the Na I RAC BEC are uncorrelated with the variations in $_\alpha$ and He I line profiles. + We conclude that we witnessed a transient phenomenon on 13th Oct. 2003 whose effects lasted a few cays at most., We conclude that we witnessed a transient phenomenon on 13th Oct. 2003 whose effects lasted a few days at most. + From the EKKepleriau dynamics of an infalline objec (Beust et al., From the Keplerian dynamics of an infalling object (Beust et al. + 2001) the most redshiftec Na I absorption component velocity (200 kun/s) should correspond. to a distance of 0.1 AU from the star., 2001) the most redshifted Na I absorption component velocity (200 km/s) should correspond to a distance of 0.1 AU from the star. + At this distance both the munber of ionizing photons and iuteusitv of the stellar wine will be too high for aux uushielded Na I to survive (Sorel et al., At this distance both the number of ionizing photons and intensity of the stellar wind will be too high for any unshielded Na I to survive (Sorelli et al. + 1996)., 1996). + Their caleulatious using phlioto-donizime codes uncer local thermodyvuauuica equilibrimu (LTE) aud non-LTE show that ouly a solid body like a comet or asteroid cau approach a hot star this close before it disintegrates and evaporates., Their calculations using photo-ionizing codes under local thermodynamical equilibrium (LTE) and non-LTE show that only a solid body like a comet or asteroid can approach a hot star this close before it disintegrates and evaporates. + Further. theoretical model caleulatious for stars up to D9 show that the size of solid bodies which are able to survive up to a distance of 0.1 AU should be at least 100 kin in diameter (Beust et al.," Further, theoretical model calculations for stars up to B9 show that the size of solid bodies which are able to survive up to a distance of 0.1 AU should be at least 100 km in diameter (Beust et al." + 2001)., 2001). + The mild Na I BEC could then be part of the falling evaporating body beiug blown away by the stellar wind or radiatiou pressure., The mild Na I BEC could then be part of the falling evaporating body being blown away by the stellar wind or radiation pressure. + However. this will also mean that the falling solid. body may have a Ligh eccentric orbit so that the blown away niaterial is not projected against the stellar surface.," However, this will also mean that the falling solid body may have a high eccentric orbit so that the blown away material is not projected against the stellar surface." +" Our Na T RAC and BEC. observatious are consistent with such an object dissociating between 2 AU and 0.1 AU from the Bh star LkIL, 231.", Our Na I RAC and BEC observations are consistent with such an object dissociating between 2 AU and 0.1 AU from the B5 star $_\alpha$ 234. +" Iu many Herbie Ac stars, variations in Na I RACs aud BAC’s usually correlates with IL, aud Ie T variations (de Winter et al."," In many Herbig Ae stars, variations in Na I RACs and BAC's usually correlates with $_\alpha$ and He I variations (de Winter et al." + 1999. Caimin et al.," 1999, Grinin et al." + 1996)., 1996). + Such Na I variations are now thought to be eenerated by maegneto-lvdrodvnamic fuuncllius of neutral gas on to the star rather than solid body iufall (Beust et al., Such Na I variations are now thought to be generated by magneto-hydrodynamic funnelling of neutral gas on to the star rather than solid body infall (Beust et al. + 2001. Mora et al.," 2001, Mora et al." + 2002. Natta et al.," 2002, Natta et al." + 2000)., 2000). +" No such correlation is present in the current LkIL,2231 data set.", No such correlation is present in the current $_\alpha$ 234 data set. +" Of particular significance is the observation on 27th Oct. 2003 when the REC and BEC of IL, axe the strongest.", Of particular significance is the observation on 27th Oct. 2003 when the REC and BEC of $_\alpha$ are the strongest. + The absence of auv Na TRAC and BEC in this observation is a strong aremucut against the Na I line variations being caused by neutral matter unnelled outo the star by magnetic fields., The absence of any Na I RAC and BEC in this observation is a strong argument against the Na I line variations being caused by neutral matter funnelled onto the star by magnetic fields. +" The variability of IL, aud We Tin LkIL,231 are due ο. episodic eas accretion. while the observed variations m Na I lines is a transicut event that exhibits the dynamical signature of a body (2100 sun in diuueter) falhug outo the star."," The variability of $_\alpha$ and He I in $_\alpha$ 234 are due to episodic gas accretion, while the observed variations in Na I lines is a transient event that exhibits the dynamical signature of a body $\geq$ 100 km in diameter) falling onto the star." + It is worth uecutionimeg here that recent studies ou meteors roni cometary origin bv Trigo-Rodriguez et al. (, It is worth mentioning here that recent studies on meteors from cometary origin by Trigo-Rodriguez et al. ( +2001) have shown greater sodimm abuudances hau those expected for iuter-planetarv dust yarticles aud. choncdiitie meteorites and Potter et al. (,2004) have shown greater sodium abundances than those expected for inter-planetary dust particles and chondritic meteorites and Potter et al. ( +2002) have detected comet ike sodium tail TOM alanet Moereurv. aud also he discovery of oeseuce of socia tail iu comets (Cremonese et al.,"2002) have detected comet like sodium tail from planet Mercury, and also the discovery of presence of sodium tail in comets (Cremonese et al." + 1997)., 1997). + Thus an infalliug solid body of asteroid size aud deficient in IT can produce the observed Na TRAC aud DEC aud no correlation with the II aud We I lines., Thus an infalling solid body of asteroid size and deficient in H can produce the observed Na I RAC and BEC and no correlation with the H and He I lines. + The dusty circtuustellar disk of the voung star, The dusty circumstellar disk of the young star +universal mechanisms (to excite or maintain warps (seeBinney1992.forareview)...,universal mechanisms to excite or maintain warps \citep[see][for a review]{binney92}. + Many of these proposed mechanisms rely on the dark halo to either stabilize warps as discrete bending modes within the halo (Sparke Casertano 1055: bul see also Binney. Jiang. Dutta 1998). or to provide the torque necessary to create the warp (Ostriker&Binney1989:Sellwood1999:Idetaetal. 2000).," Many of these proposed mechanisms rely on the dark halo to either stabilize warps as discrete bending modes within the halo (Sparke Casertano 1988; but see also Binney, Jiang, Dutta 1998), or to provide the torque necessary to create the warp \citep{ostriker and binney89,debattista and sellwood99,ideta et al00}." +. Other proposed mechanisms include the inlall of intergalactic gas (Lóppez-Corredoira. Belancort-Rijo. Beckman 2002a). magnetic fields (Dattaner. Florido. Sanchez-Saavedra 1990). and interactions with satellite galaxies (e.g..Tang&Carlberg1997).," Other proposed mechanisms include the infall of intergalactic gas (Lóppez-Corredoira, Betancort-Rijo, Beckman 2002a), magnetic fields (Battaner, Florido, Sanchez-Saavedra 1990), and interactions with satellite galaxies \citep[e.g.,][]{huang and carlberg97}." +. Each of these mechanisms can. in particular circumstances. produce realistic-looking ealactie warps.," Each of these mechanisms can, in particular circumstances, produce realistic-looking galactic warps." + Although no single mechanism appears universal enough to account for all warps. the evolution toward a bending mode (even when no discrete mode exists) appears enough like an observed warp (IllolÀer&Sparke1994) that warping may be a generic response of disks to the individual perturbations thev experience.," Although no single mechanism appears universal enough to account for all warps, the evolution toward a bending mode (even when no discrete mode exists) appears enough like an observed warp \citep{hofner and sparke94} that warping may be a generic response of disks to the individual perturbations they experience." + In this case. we should look at individual warped galaxies lor specific evidence of particular perturbations (hat explain their warps rather (han search for a universal mechanism that may not exist.," In this case, we should look at individual warped galaxies for specific evidence of particular perturbations that explain their warps rather than search for a universal mechanism that may not exist." + The Magellanic Clouds have been proposed as the perturbation responsible for the Milky Ways warp., The Magellanic Clouds have been proposed as the perturbation responsible for the Milky Way's warp. + While IIunter&Toomre(1969). found that the tidal distortion from the clouds alone is not sufficient to cause the observed warp. Weinberg(1998) proposed that orbiting satellites could set up wakes in (he Milky Ways halo which could provide the necessary (torque.," While \citet{hunter and toomre69} found that the tidal distortion from the clouds alone is not sufficient to cause the observed warp, \citet{weinberg98} proposed that orbiting satellites could set up wakes in the Milky Way's halo which could provide the necessary torque." + Tsuchiva(2002) performed self-consistent simulations of such a system and confirmed that for a sufficiently massive halo (2.1x107.M. ). the magnitude of the torque can be increased enough (o cause a warp of the same magnitude as the Milkv Ways.," \citet{tsuchiya02} performed self-consistent simulations of such a system and confirmed that for a sufficiently massive halo $2.1\times10^{12}~M_{\Sun}$ ), the magnitude of the torque can be increased enough to cause a warp of the same magnitude as the Milky Way's." + The Alagellanie Clouds orbit about the center of the Galaxy in a direction orthogonal to the line of nodes. ie. near the line of maxim warp (see Figure 1)).," The Magellanic Clouds orbit about the center of the Galaxy in a direction orthogonal to the line of nodes, i.e., near the line of maximum warp (see Figure \ref{schematic}) )." +" Garciaa-Rauiz. κο, Dubinski (2000) demonstrated (hat the warp caused by a satellite will have its line of nodes oriented the satellite's orbit."," Garcíaa-Ruiz, Kuijken, Dubinski (2000) demonstrated that the warp caused by a satellite will have its line of nodes oriented the satellite's orbit." + A simple wav of understanding this result is {ο recognize (hat a torque is a transfer of angular momentum. and therefore the disk will acquire angular momentum along the same axis as (he orbital angular momentum ol the satellite which is providing the torque. and (lt toward that axis.," A simple way of understanding this result is to recognize that a torque is a transfer of angular momentum, and therefore the disk will acquire angular momentum along the same axis as the orbital angular momentum of the satellite which is providing the torque, and tilt toward that axis." + Therefore. the Magellanic Clouds are a bad candidate for producing the Milky Way warp.," Therefore, the Magellanic Clouds are a bad candidate for producing the Milky Way warp." + The orbital plane of the Sagittarius dwarl galaxy. (bata. Gilmore. Irwin 1994) does intersect the line of nodes. suggesting (hat it may be a good candidate for producing the Milky Wav warp (Lin1996).," The orbital plane of the Sagittarius dwarf galaxy (Ibata, Gilmore, Irwin 1994) does intersect the line of nodes, suggesting that it may be a good candidate for producing the Milky Way warp \citep{lin96}." +. It is located behind the Galactic bulge and is on a nearly polar orbit (Ibataοἱal., It is located behind the Galactic bulge and is on a nearly polar orbit \citep{ibata et al97}. +1997).. Ibata&BRazoumov(1998). performed simulations which suggest that the passage of a sufficiently massive Ser (5x10? AZ.) through the disk could produce a warp., \citet{ibata and razoumov98} performed simulations which suggest that the passage of a sufficiently massive Sgr $5\times 10^9~M_{\Sun}$ ) through the disk could produce a warp. + Alternatively. its gravitational tides or the tides of a wake it produces in the dark," Alternatively, its gravitational tides or the tides of a wake it produces in the dark" +"angular size of a speckle is 6,5444,=A/D, where A is the wavelength of observation and D is the telescope diameter.","angular size of a speckle is $\theta_{speckle}=\lambda/D$, where $\lambda$ is the wavelength of observation and $D$ is the telescope diameter." + The planet is assumed to be an unresolved point source and its image has a full width at half-maximum (FWHM) that is approximately equal to the diffraction limit., The planet is assumed to be an unresolved point source and its image has a full width at half-maximum (FWHM) that is approximately equal to the diffraction limit. +" Differentiating a faint planet point-spread function (PSF) from a speckle is therefore impossible with a single image, and without using a temporal sequence or other spectral information."," Differentiating a faint planet point-spread function (PSF) from a speckle is therefore impossible with a single image, and without using a temporal sequence or other spectral information." +" The position and size of a speckle in the field of view is wavelength-dependent: as the wavelength increases from A, to 4», the FWHM of a single speckle and its angular separation from the star both increase by a factor 45/4."," The position and size of a speckle in the field of view is wavelength-dependent: as the wavelength increases from $\lambda_{1}$ to $\lambda_{2}$, the FWHM of a single speckle and its angular separation from the star both increase by a factor $\lambda_{2}/\lambda_{1}$." +" Following the wavelength axis at a fixed angular separation from the star, one would measure a strong modulation in the intensity as the Airy rings and speckles cross this position (see?,Fig.24).."," Following the wavelength axis at a fixed angular separation from the star, one would measure a strong modulation in the intensity as the Airy rings and speckles cross this position \citep[see][Fig.~24]{sparks2002}." + A possible faint companion signal would remain undetected in the modulations due to the wavelength dependence of the speckle pattern., A possible faint companion signal would remain undetected in the modulations due to the wavelength dependence of the speckle pattern. +" However, a fixed physical object (e.g. a planetary companion) will not change its position with wavelength: only its FWHM will be multiplied by the aforementioned factor."," However, a fixed physical object (e.g. a planetary companion) will not change its position with wavelength: only its FWHM will be multiplied by the aforementioned factor." +" This property enables a good determination and subtraction of wavelength dependent features at a given position, and allows the detection of physical objects in the vicinity of the star."," This property enables a good determination and subtraction of wavelength dependent features at a given position, and allows the detection of physical objects in the vicinity of the star." +" In the case of long slit spectroscopy, an infinity of images of the slit at increasing wavelengths are dispersed and superimposed to create a spectrum such as that shown in Fig. [I]."," In the case of long slit spectroscopy, an infinity of images of the slit at increasing wavelengths are dispersed and superimposed to create a spectrum such as that shown in Fig. \ref{figure:typical_spectrum}." + The wavelength-dependent artifacts such as speckles create oblique in the spectrum., The wavelength-dependent artifacts such as speckles create oblique in the spectrum. +" A bright speckle will see its size increasedlined] with wavelength, creating a bright line in the spectrum with increasing width."," A bright speckle will see its size increased with wavelength, creating a bright line in the spectrum with increasing width." +" A physical object creates a straight line at a fixed radial separation from the star, and its FWHM increases linearly with wavelength."," A physical object creates a straight line at a fixed radial separation from the star, and its FWHM increases linearly with wavelength." +" Our method uses these geometrical properties to separate the planet and the star spectra, as proposed by ? with spectral deconvolution (SD) in the case of IFS observations."," Our method uses these geometrical properties to separate the planet and the star spectra, as proposed by \citet{thatte2007} with spectral deconvolution (SD) in the case of IFS observations." +" In the image representing the spectrum, each column of pixels corresponds to a spectral interval AA of a few nanometers, centered on a particular wavelength."," In the image representing the spectrum, each column of pixels corresponds to a spectral interval $\Delta\lambda$ of a few nanometers, centered on a particular wavelength." +" From now on, we assume that a given column of pixel i in the spectrum corresponds to a single wavelength equal to the central wavelength 4; of the spectral interval ΔΑ."," From now on, we assume that a given column of pixel $i$ in the spectrum corresponds to a single wavelength equal to the central wavelength $\lambda_{i}$ of the spectral interval $\Delta\lambda_{i}$." +" This approximation is valid as long as AA;/A;«1, which is the case for the wavelength range and spectral interval considered in our work (see Sect. Bp "," This approximation is valid as long as $\Delta\lambda_{i}/\lambda_{i} \ll 1$, which is the case for the wavelength range and spectral interval considered in our work (see Sect. \ref{section:lss_simulations}) )" +where ΔΑΙ/Aj«0.015 in any case., where $\Delta\lambda_{i}/\lambda_{i} < 0.015$ in any case. +" The first step corrects the spectral dependency of the speckles by spatially rescaling each column of pixels by a factor αι=Apo/dA;, where Αρ is the shortest spectrum wavelength."," The first step corrects the spectral dependency of the speckles by spatially rescaling each column of pixels by a factor $\alpha_{i}=\lambda_{0}/\lambda_{i}$, where $\lambda_{0}$ is the shortest spectrum wavelength." +" Since o; is not an integer for most columns, each column was resampled in the spatial direction to rescale its length from one integer number of pixels to another integer number of pixels."," Since $\alpha_{i}$ is not an integer for most columns, each column was resampled in the spatial direction to rescale its length from one integer number of pixels to another integer number of pixels." +" The number of resolution elements is fixed by the instrumental setup (i.e. the platescale) in the spatial direction, and by the dispersive element in the spectral direction."," The number of resolution elements is fixed by the instrumental setup (i.e. the platescale) in the spatial direction, and by the dispersive element in the spectral direction." +" In fact, the spectrum is resampled within a grid ~15 times smaller than the original pixel grid before rescaling the columns by the factor αι."," In fact, the spectrum is resampled within a grid $\sim$ 15 times smaller than the original pixel grid before rescaling the columns by the factor $\alpha_{i}$." +" When calculating the new integer number of pixels for column i, this spatial oversampling allows us to reduce the round-off error introduced by the non-integer factor αι."," When calculating the new integer number of pixels for column $i$, this spatial oversampling allows us to reduce the round-off error introduced by the non-integer factor $\alpha_{i}$." +" With the approximation mentioned in the previous paragraph, neglecting the spectral interval A4; translates into an error of less than in αι."," With the approximation mentioned in the previous paragraph, neglecting the spectral interval $\Delta\lambda_{i}$ translates into an error of less than in $\alpha_{i}$." + The final rescaled spectrum is shown in Fig.[]., The final rescaled spectrum is shown in Fig. \ref{figure:rescaled_spectrum}. +" The rescaling corrects the spectral dependence of each speckle, ensuring that they follow horizontal straight lines, while the companion spectrum follows a 1/4 law."," The rescaling corrects the spectral dependence of each speckle, ensuring that they follow horizontal straight lines, while the companion spectrum follows a $1/\lambda$ law." + Speckles are induced by random phase errors in the incoming wavefront that scatters the starlight., Speckles are induced by random phase errors in the incoming wavefront that scatters the starlight. +" The fraction of light scattered in each speckle coming from the companion is negligible compared to the scattered light of the star, so it is safe to assume"," The fraction of light scattered in each speckle coming from the companion is negligible compared to the scattered light of the star, so it is safe to assume" +αἱ lo for H(z)--DAO--CMDB.,at $1\sigma$ for $H(z)$ +BAO+CMB. + However. obviously. with the two joint analyses including the SNe Ia data. ACDAL is still included within 1o error region.," However, obviously, with the two joint analyses including the SNe Ia data, $\Lambda$ CDM is still included within $1\sigma$ error region." + Moreover. comparing Fig.," Moreover, comparing Fig." + Ib to Fig., 1b to Fig. + 1e. we can find the confidence regions of (/(2)+BAO+CAIB cata are in good agreement with that of SNe Ia--DAO--CMD data: this situation has also been noted when constraining on the ACDM. XCDM scenario (Zhaietal.2010). and the interacting dark matter models without dark energv (Cao.Zhu&Liang2011)..," 1c, we can find the confidence regions of $H(z)$ +BAO+CMB data are in good agreement with that of SNe Ia+BAO+CMB data; this situation has also been noted when constraining on the $\Lambda$ CDM, XCDM scenario \citep{Zhai10} and the interacting dark matter models without dark energy \citep{Cao11a}." + If the interaction term is Q=3254// px. in spatially flat FRW metric. for the the ο IDE model with a constant Eos of dark energy τῶν. the Friedinann equation is The joint confidence regions in i6 y-54 plane wilh different observational data sets (Jf(2). 1(2)4+BAO+CA3IB. SNe lat+¢BAO+CAB. aad. H(z)2-SNe InDAO--CMDJ) for the 5; IDE model are showedin Fig.," If the interaction term is $Q=3\gamma_d H\rho_X$ , in spatially flat FRW metric, for the the $\gamma_d$ IDE model with a constant EoS of dark energy $w_X$, the Friedmann equation is The joint confidence regions in $w_X$ $\gamma_d$ plane with different observational data sets $H(z)$, $H(z)$ +BAO+CMB, SNe Ia+BAO+CMB, and $H(z)$ +SNe Ia+BAO+CMB) for the $\gamma_d$ IDE model are showed in Fig." + 2., 2. + We also present the best-fit values of parameters will 1-0 and 2-0 uncertainties [or the5; IDE modelin Table 1.., We also present the best-fit values of parameters with $\sigma$ and $\sigma$ uncertainties for the $\gamma_d$ IDE model in Table \ref{tab2}. + With the /(z .oe only (Fig., With the $H(z)$ data only (Fig. + 2a). the best-fit values of (he parameters (ey. ty) ave wy=—130 andsy=M57.," 2a), the best-fit values of the parameters $w_X, \gamma_d$ ) are $w_X=-1.30$ and $\gamma_d=-0.057$." + With /7(2)4+-BAO+CAIB (Fig., With $H(z)$ +BAO+CMB (Fig. + 2b). the best-fit values at 1-7 confidence level are (τῶν.=(—124iris.—0.040.i3).," 2b), the best-fit values at $\sigma$ confidence level are $(w_X, +\gamma_d)=(-1.24_{-0.19}^{+0.17}, -0.040_{-0.042}^{+0.042})$." + Fitting results [rom the joint data with SNe Ia4-DAO--CMD are given in Fig., Fitting results from the joint data with SNe Ia+BAO+CMB are given in Fig. + 2e with the best-lit values ο=(71.02.tras.20.023.min).," 2c with the best-fit values $(w_X, +\gamma_d)=(-1.02_{-0.09}^{+0.09},-0.023_{-0.040}^{+0.039})$." + Though the special case of the ACDM with no interaction is excluded at lo for /7€:)2-DAO--CMD., Though the special case of the $\Lambda$ CDM with no interaction is excluded at $1\sigma$ for $H(z)$ +BAO+CMB. + Fitting results from the joint data with //(z)-2-9Ne Ia-4-DAO--CMD given in Fig., Fitting results from the joint data with $H(z)$ +SNe Ia+BAO+CMB given in Fig. + 2d with the best-fit values —0.030.082) indicate that ACDM is still included within lo error region.," 2d with the best-fit values $(w_X, \gamma_d)=(-1.10_{-0.13}^{+0.13},-0.030_{-0.037}^{+0.037})$ indicate that $\Lambda$ CDM is still included within $1\sigma$ error region." + From Fig., From Fig. + 1-2 and Table 1.. it is obvious that the constraints on both interacting scenarios favor 5«0. which indicates that the energy is transferred. [rom dark matter to dark energy and the coincidence problem is «uite severe. a result consistent wilh the previous results by using other observational data including the 182 Gold SNe Ia and 397 Constitution 5Ne la samples (Chenetal.2010:Feng2007).," 1-2 and Table \ref{tab2}, it is obvious that the constraints on both interacting scenarios favor $\gamma<0$, which indicates that the energy is transferred from dark matter to dark energy and the coincidence problem is quite severe, a result consistent with the previous results by using other observational data including the 182 Gold SNe Ia and 397 Constitution SNe Ia samples \citep{Chen10,Feng07}." +. In addition. the constraining; results in this work with the joint observational data including the //(2) data are more stringent than (he previous results lor constraining the interaction term with othercombined observations arising from the 182 Gold SNe Ia samples. the shilt parameter of CMD given by the 2-vear WALAP observations. the BAO measurement [rom SDSS and age estimates of 35 galaxies (Fengetal. 2003)..," In addition, the constraining results in this work with the joint observational data including the $H(z)$ data are more stringent than the previous results for constraining the interaction term with othercombined observations arising from the 182 Gold SNe Ia samples, the shift parameter of CMB given by the 3-year WMAP observations, the BAO measurement from SDSS and age estimates of 35 galaxies \citep{Feng08}. ." +Galaxyv-ealaxy interactions/nmergers are very important in (he cosmic evolution of ealaxies.,Galaxy-galaxy interactions/mergers are very important in the cosmic evolution of galaxies. + In the hierarchical galaxy formation paradigm. galaxies ancl galactic structures are," In the hierarchical galaxy formation paradigm, galaxies and galactic structures are" +combining the ten per cent of the images with the best seeing.,combining the ten per cent of the images with the best seeing. + This reference image was then convolved with a spatially varying kernel to degrade its seeing to match the seeing of each individual image in turn., This reference image was then convolved with a spatially varying kernel to degrade its seeing to match the seeing of each individual image in turn. +" Each individual image was then subtracted from the appropriately convolved reference image, with the result of creating a subtracted image containing the residuals (the outbursting CV is obvious in these subtracted images — see Fig."," Each individual image was then subtracted from the appropriately convolved reference image, with the result of creating a subtracted image containing the residuals (the outbursting CV is obvious in these subtracted images – see Fig." + 2 (c))., \ref{fig:m22smarts}~ (c)). +" All the subtracted images were median-combined to produce an image,var.fits, showing — in theory — only the variable stars, which appear as stellar-like profiles."," All the subtracted images were median-combined to produce an image, showing – in theory – only the variable stars, which appear as stellar-like profiles." +" In practice, saturated stars and cosmic rays lead to many spurious detections."," In practice, saturated stars and cosmic rays lead to many spurious detections." + The light curve of the CV was extracted from the subtracted images in theV andB bands using the routine packaged with and plotted in units of differential flux relative to the reference image in which the CV was in quiescence., The light curve of the CV was extracted from the subtracted images in the and bands using the routine packaged with and plotted in units of differential flux relative to the reference image in which the CV was in quiescence. + We retrieved archival multicolour images of M22 from the Multimission Archive at A list of the observations is given in Table 1.., We retrieved archival multicolour images of M22 from the Multimission Archive at A list of the observations is given in Table \ref{table:obs}. +" CV1 appears on chip WFA of the Wide-Field Planetary Camera 2 (WFPC2), as seen in Fig. 3::"," CV1 appears on chip WF4 of the Wide-Field Planetary Camera 2 (WFPC2), as seen in Fig. \ref{fig:m22find}:" +" its coordinates are a=18°36™24°.66, 6=—23?54'35""5 (J2000;Sahuetal.2001)."," its coordinates are $\alpha= 18^{\mathrm{h}}36^{\mathrm{m}}24^{\mathrm{s}}\!.66$, $\delta= -23\degr54\arcmin35\arcsec\!.5$ \citep[J2000;][]{sahu01}." +". In order to correct the counts in each pixel of the pipeline-calibrated images for the effects of geometric distortion (which appears at the 1-2 per cent level near the edges of the images and up to a maximum of 4-5 per cent in the corners, for fixed-aperture photometry), we multiplied each image by the correction image obtained from the archive."," In order to correct the counts in each pixel of the pipeline-calibrated images for the effects of geometric distortion (which appears at the 1–2 per cent level near the edges of the images and up to a maximum of 4–5 per cent in the corners, for fixed-aperture photometry), we multiplied each image by the correction image obtained from the archive." +" Cosmic rays were removed in with the tool for the co-aligned images taken with theF255W filter (nUVoss-band) and with the Laplacian edge-detection routine (vanDokkum2001) for the individual images in theF336W and filters (U336- and Vsss-band, respectively)."," Cosmic rays were removed in with the tool for the co-aligned images taken with the filter $nUV_{255}$ -band) and with the Laplacian edge-detection routine \citep{dokkum01} for the individual images in the and filters $U_{336}$ - and $V_{555}$ -band, respectively)." +" We performed small-aperture photometry in the U336 and Vsss bands with the (Stetson1987) package in on all objects detected — including CV1 - in the V555-band image on the PC, WF2 and WF4 chips."," We performed small-aperture photometry in the $U_{336}$ and $V_{555}$ bands with the \citep{stetson87} package in on all objects detected – including CV1 – in the $V_{555}$ -band image on the PC, WF2 and WF4 chips." + We corrected the results to the aperture for which the zeropoint was defined., We corrected the results to the aperture for which the zeropoint was defined. + The magnitudes were calibrated on the STMAG flux-based system using the zeropoint calculated from the PHOTFLAM header keyword., The magnitudes were calibrated on the STMAG flux-based system using the zeropoint calculated from the PHOTFLAM header keyword. + All the magnitudes were corrected for charge transfer efficiency (CTE) losses, All the magnitudes were corrected for charge transfer efficiency (CTE) losses + (Waeneretal.1991: I&ougetal.20023). (Waeneretal.1992: Casaresetal.1993:: Pavleukoetal.1996:: TIvuesetal. 2002: Zurita.Casares.&Shahbaz2003:: Shahbazetal. 200511. IR (Samwaletal.1996).. aud radio (Ποήπιοetal.2000).. but none of these studies were coordinated.," \citealt{Wagner:1994a}; \citealt{Kong:2002a}) \citealt{Wagner:1992a}; \citealt{Casares:1993a}; \citealt{Pavlenko:1996a}; \citealt{Hynes:2002a}; ; \citealt{Zurita:2003a}; \citealt{Shahbaz:2004a}) ), IR \citep{Sanwal:1996a}, and radio \citep{Hjellming:2000a}, but none of these studies were coordinated." + ντοςetal.(2002). established that optical eniission line variations are correlated with the optical continu., \citet{Hynes:2002a} established that optical emission line variations are correlated with the optical continuum. + They also found that the eiissiou line flares exhibited a double-peaked line profile. suggestive of enuüssion distributed across the accretion disk (see e.g. Tome&Alarsh 1986)) rather thin ariiug iu localized regious.," They also found that the emission line flares exhibited a double-peaked line profile, suggestive of emission distributed across the accretion disk (see e.g., \citealt{Horne:1986a}) ) rather than arising in localized regions." + This was attributed to irradiation of the outer disk bv the variable N-rav source. and hence it was predieted that the N-rav variations should be correlated with the optical.," This was attributed to irradiation of the outer disk by the variable X-ray source, and hence it was predicted that the X-ray variations should be correlated with the optical." + Such correlated variability. also attributed to mradiation. is commonly seen d X-rav bright states in both neutron star svstenis and black holes citealtCadudlau:1978v: Petroetal. 1981: IIxnesetal. 1998: and mauv other works). but had not been directly observed iu quiescent systems.," Such correlated variability, also attributed to irradiation, is commonly seen in X-ray bright states in both neutron star systems and black holes \\citealt{Grindlay:1978a}; \citealt{Petro:1981a}; \citealt{Hynes:1998a}; and many other works), but had not been directly observed in quiescent systems." + It is also usually oulv detected in the optical contimmmur viclding no kincmatic information., It is also usually only detected in the optical continuum yielding no kinematic information. + ere we report initial results from a coordinated. multiwavelength campaign to test this prediction.," Here we report initial results from a coordinated, multiwavelength campaign to test this prediction." + This included X-ray. ucar-UV. optical. aud radio coverage. but this letter discusses oulv the results from comparing N-rav data with optical spectroscopy.," This included X-ray, near-UV, optical, and radio coverage, but this letter discusses only the results from comparing X-ray data with optical spectroscopy." + Future works will study the variability properties m mere detail aud exanune the broad-band spectral energy distribution., Future works will study the variability properties in more detail and examine the broad-band spectral energy distribution. + oobservations oi 2003 July 28/29 used the ACTS camera. in a sinele kks observation spaunius binary phases 0.510.62.," observations on 2003 July 28/29 used the ACIS camera, in a single ks observation spanning binary phases 0.51–0.62." + The source was positioned ou the ACIS-S3 chip aud the 1/8 sub-uray mode was used to reduce the frune-tinme to 0.[ss aud hence eusure that pile-up was neeheible., The source was positioned on the ACIS-S3 chip and the $1/8$ sub-array mode was used to reduce the frame-time to s and hence ensure that pile-up was negligible. + Data analysis used. Ciao 3.0., Data analysis used Ciao 3.0. + Source events were extracted from a radius aperture. retaining events withenergies of kkoV: a total of 1911 such eveuts were recorded," Source events were extracted from a radius aperture, retaining events withenergies of keV; a total of 1941 such events were recorded" + , +llaving presented our results on upper limits to the anisotropy [actor for turbulence in the Local Clouds. we now consider the corresponding quantity in the solar wind. which has the benefit of direct. in-situ measurements.,"Having presented our results on upper limits to the anisotropy factor for turbulence in the Local Clouds, we now consider the corresponding quantity in the solar wind, which has the benefit of direct, in-situ measurements." + SpanglerandSpitller.(2004) resolved solar wind magnetic field fluctuations into components parallel aud. perpendicular to the inlerplanelarv magnetic field., \cite{Spangler04} resolved solar wind magnetic field fluctuations into components parallel and perpendicular to the large-scale interplanetary magnetic field. + The data came from the magnetometer of the WIND spacecraft at a heliocentric distance of about 1 AU., The data came from the magnetometer of the WIND spacecraft at a heliocentric distance of about 1 AU. + SpanglerandSpitler(2004) sed 66 intervals of one hour duration during slow solar wind conditions. ancl report their results in terms of modulation indices my) and my of fluctuations parallel and perpendicular to the mean field!.. Where 06) and 0b. ave the rms values of the ΕΕ in the magnetic field components parallel ancl perpendicular. respectively. to the mean field By.," \cite{Spangler04} used 66 intervals of one hour duration during slow solar wind conditions, and report their results in terms of modulation indices $m_{B \parallel}$ and $m_{B \perp}$ of fluctuations parallel and perpendicular to the mean field, where $\delta b_{\parallel}$ and $\delta b_{\perp}$ are the rms values of the fluctuations in the magnetic field components parallel and perpendicular, respectively, to the mean field $B_0$." + report mean values [or my) and my of 0.0321 and 0.112. respectively.," \cite{Spangler04} report mean values for $m_{B \parallel}$ and $m_{B \perp}$ of 0.0321 and 0.112, respectively." + The means are for the distribution of values measured in the 66 data intervals., The means are for the distribution of values measured in the 66 data intervals. +" As noted in (2004).. my, should be larger than my) because 1t possesses contributions [rom (wo turbulent field components rather (han just one."," As noted in \cite{Spangler04}, $m_{B \perp}$ should be larger than $m_{B \parallel}$ because it possesses contributions from two turbulent field components rather than just one." + The degree of anisotropy can be determined by comparing my) to mp/v2. Spa, The degree of anisotropy can be determined by comparing $m_{B \parallel}$ to $m_{B \perp}/\sqrt 2$. +nglerandSpitler(2004) found that the turbulent (Inctuations in their study. were Alfvénnic. in the sense that ryde=abBy 5ο the measured magnetic field modulation indices may be considered proxies for modulation indices of the velocity. Huctuations.," \cite{Spangler04} found that the turbulent fluctuations in their study were Alfvénnic, in the sense that $\frac{\delta v}{V_A} = \frac{\delta b}{B_0}$, so the measured magnetic field modulation indices may be considered proxies for modulation indices of the velocity fluctuations." + With (his assumption. we have where in Equations (3) and (9) we make the connection between the root-mean-square velocity fluctuations and the velocity scales Vj and V. of the fluctuation distribution function," With this assumption, we have where in Equations (8) and (9) we make the connection between the root-mean-square velocity fluctuations and the velocity scales $V_{\parallel}$ and $V_{\perp}$ of the fluctuation distribution function" +"UV-Visual colour-magnitude diagrams of the two very massive globular clusters. co Cen and NGC 2808. show a rather puzzling ""hook-like feature at the hot end of their extended horizontal branches with stars lying below the canonical horizontal branch (Whitney et al. 1998::","UV-Visual colour-magnitude diagrams of the two very massive globular clusters, $\omega$ Cen and NGC 2808, show a rather puzzling “hook-like” feature at the hot end of their extended horizontal branches with stars lying below the canonical horizontal branch (Whitney et al. \cite{whro98};" + D'Cruz et al. 2000:;, D'Cruz et al. \cite{dcoc00}; + Brown et al. 2001))., Brown et al. \cite{brsw01}) ). + These stars cannot be explained within the framework of canonical stellar evolution., These stars cannot be explained within the framework of canonical stellar evolution. +" Brown et ((2001)) have proposed a ""flash-mixing"". scenario to. explain the blue hook stars.", Brown et \cite{brsw01}) ) have proposed a “flash-mixing” scenario to explain the blue hook stars. + According to this scenario stars which lose an unusually large amount of mass will leave the red giant branch (RGB) before the helium flash and will move quickly to the (helium-core) white dwarf cooling curve before igniting helium (Castellani Castellani 1993:; D'Cruz et citededo96:; Brown et al. 2001)., According to this scenario stars which lose an unusually large amount of mass will leave the red giant branch ) before the helium flash and will move quickly to the (helium-core) white dwarf cooling curve before igniting helium (Castellani Castellani \cite{caca93}; D'Cruz et \\cite{dcdo96}; Brown et al. \cite{brsw01}) ). +" However. the evolution of these ""late hot helium flashers"" differs dramatically from the evolution of stars which undergo the helium flash on the RGB."," However, the evolution of these “late hot helium flashers” differs dramatically from the evolution of stars which undergo the helium flash on the RGB." + Ordinarily when a star flashes at the tip of the RGB or shortly thereafter. the large entropy barrier of its strong hydrogen-burning shell prevents the products of helium burning from being mixed to the surface.," Ordinarily when a star flashes at the tip of the RGB or shortly thereafter, the large entropy barrier of its strong hydrogen-burning shell prevents the products of helium burning from being mixed to the surface." + Such canonical stars will evolve to the zero-age horizontal branch ZAHB) without any change in their hydrogen-rich envelope composition., Such canonical stars will evolve to the zero-age horizontal branch ) without any change in their hydrogen-rich envelope composition. + In contrast. stars that ignite helium on the white dwarf cooling curve. where the hydrogen-buning shell is much weaker. will undergo extensive mixing between the helium- and carbon-rich. core and the hydrogen envelope (Sweigart 1997:: Brown et al. 2001," In contrast, stars that ignite helium on the white dwarf cooling curve, where the hydrogen-burning shell is much weaker, will undergo extensive mixing between the helium- and carbon-rich core and the hydrogen envelope (Sweigart \cite{swei97}; Brown et al. \cite{brsw01};" + Cassist et al. 2003))., Cassisi et al. \cite{cas03}) ). +" Depending on where the helium flash occurs along the white dwarf cooling curve. the envelope hydrogen will be mixed either deeply into the core (""deep mixing’) or only with a convective shell in the outer part of the core (""shallow mixing”)."," Depending on where the helium flash occurs along the white dwarf cooling curve, the envelope hydrogen will be mixed either deeply into the core (“deep mixing”) or only with a convective shell in the outer part of the core (“shallow mixing”)." + In the case of deep mixing virtually all of the envelope hydrogen ts burned while in shallow mixing some of the envelope hydrogen remains after the mixing phase (Lanz et al. 2004))., In the case of deep mixing virtually all of the envelope hydrogen is burned while in shallow mixing some of the envelope hydrogen remains after the mixing phase (Lanz et al. \cite{labr04}) ). + One of the most robust predictions of the flash-mixing scenario Is an increase in the surface abundance of carbon to - (deep mixing) or (shallow mixing) by mass., One of the most robust predictions of the flash-mixing scenario is an increase in the surface abundance of carbon to - (deep mixing) or (shallow mixing) by mass. + This increase ts set by the carbon production during the helium flash and is nearly independent of the stellar parameters., This increase is set by the carbon production during the helium flash and is nearly independent of the stellar parameters. + Nitrogen may also be enhanced due to the burning of hydrogen on triple- carbon during the flash-mixing phase., Nitrogen may also be enhanced due to the burning of hydrogen on $\alpha$ carbon during the flash-mixing phase. + For both deep and shallow mixing. the blue hook stars should be helium-rich comparedto the canonical EHB stars.," For both deep and shallow mixing, the blue hook stars should be helium-rich comparedto the canonical EHB stars." + Alternatively. the recently observed split among the main sequence stars of co Cen and 22808 (Piotto et al. 2005. 2007))," Alternatively, the recently observed split among the main sequence stars of $\omega$ Cen and 2808 (Piotto et al. \cite{pivi05,pibe07}) )" + has been attributed to a sub-population of stars with helium abundances as large as Y «0.4 (Norris 2004:: D'Antona et al. 2005::, has been attributed to a sub-population of stars with helium abundances as large as $Y$$\approx$ 0.4 (Norris \cite{norr04}; D'Antona et al. \cite{dabe05}; + D'Antona Ventura 2007:; see Newsham Terndrup 2007 for cautionary remarks)., D'Antona Ventura \cite{dave07}; see Newsham Terndrup \cite{nete07} for cautionary remarks). + Lee et al. (2005)), Lee et al. \cite{lejo05}) ) + have suggested that the blue hook stars are the progey of these proposed helium-rich main sequence stars., have suggested that the blue hook stars are the progeny of these proposed helium-rich main sequence stars. + If the blue hook stars were to be explained by the helium-enrichmet scenario. their helium abundance should not exceed ¥=0.4 and carbon should not be enriched at all.," If the blue hook stars were to be explained by the helium-enrichment scenario, their helium abundance should not exceed $Y$$\approx$ 0.4 and carbon should not be enriched at all." + Spectroscopic observations of the blue (and supposedly heltum-rich) main sequence stars in ω Cen yield a carbon abundance of [C/M] = 0.0 (Piotto et al. 2005))., Spectroscopic observations of the blue (and supposedly helium-rich) main sequence stars in $\omega$ Cen yield a carbon abundance of [C/M] = 0.0 (Piotto et al. \cite{pivi05}) ). + This carbon abundance will decrease further as the stars ascend the red giant branch. due to the extra-mixing process that occurs in metal-poor red giants (Gratton et al. 20000:," This carbon abundance will decrease further as the stars ascend the red giant branch, due to the extra-mixing process that occurs in metal-poor red giants (Gratton et al. \cite{grsn00};" + Kraft 1994))., Kraft \cite{kraf94}) ). + Origlia et al. (2003)), Origlia et al. \cite{orfe03}) ) + have confirmed that the RGB stars in w Cen have the low C/UC ratios (x4) and low average carbon abundances ({C/Fe] = —0.2) expected from this extra-mixing., have confirmed that the RGB stars in $\omega$ Cen have the low $^{12}$ $^{13}$ C ratios $\approx$ 4) and low average carbon abundances ([C/Fe] = $-$ 0.2) expected from this extra-mixing. + Thus the helrum-enrichment scenario predicts a carbon abundance by mass in the blue hook stars of less than 0.1%.. Le. at least a factor of 10 smaller than the carbon abundance predicted by the flash-mixing scenario.," Thus the helium-enrichment scenario predicts a carbon abundance by mass in the blue hook stars of less than , i.e., at least a factor of 10 smaller than the carbon abundance predicted by the flash-mixing scenario." + Previous spectra of the blue hook stars in «o Cen (Moehler et al. 2002)), Previous spectra of the blue hook stars in $\omega$ Cen (Moehler et al. \cite{mosw02}) ) + and NGC 2808 (Moehler et, and NGC 2808 (Moehler et +llowever. as reviewed by (Stetson.Bruntt&Grundhal2003:Carney.2005).. chemical abundances are difficult to measure in this moderately distant aid reddened cluster.,"However, as reviewed by \citep{stet03,carney05}, chemical abundances are difficult to measure in this moderately distant and reddened cluster." + Hs most luminous stars are relatively faint and the combination of low effective temperature and high metallicity make hieh resolution optical spectra difficult to analyze., Its most luminous stars are relatively faint and the combination of low effective temperature and high metallicity make high resolution optical spectra difficult to analyze. + Thus. metallicity estimates have been derived mainly using three different. technique such as photometric metallicity indicators (Janes1984:Canternaetal.19386): low-- and moderate-resolution spectroscopy (Friel&Janes1993:PetersonGreenWorthev&Jowell 2003):: and model isochrones (Stetson.Druntt&Grundhbal2003:Carney.Lee&Dodson2005.andreference therein).," Thus, metallicity estimates have been derived mainly using three different technique such as photometric metallicity indicators \citep{j84,cat86}; - and -resolution spectroscopy \citep{fj93,pet98,friel02,wj03}; ; and model isochrones \citep[][and reference therein]{stet03,carney05}." +. To summarize. (he metallicity proposed for this eluster is in the range +0.11+0.44 dex.," To summarize, the metallicity proposed for this cluster is in the range +0.11–+0.44 dex." + The first work at mecdium-high resolution is (he one by Peterson&Green(1998).. who measured a sample of warm HB stars αἱ R=20.000. finding an iron abundance —-4-0.42:0.1. a modest. (if any) a-enhancement (within a factor of 2). and about solar |C'/Fe].," The first work at medium-high resolution is the one by \citet{pet98}, who measured a sample of warm HB stars at R=20,000, finding an iron abundance $\pm$ 0.1, a modest (if any) $\alpha$ -enhancement (within a factor of 2), and about solar [C/Fe]." + Very recently. (wo other spectroscopic studies on clump and Rec Giant Branch (RGB) stars. have been performed by. Carraroetal.(2006) and Grattonetal.(2006) finding [Fe/I]J=+0.3940.01 and |Fe/1I]—--0.472:0.04. respectively.," Very recently, two other spectroscopic studies on clump and Red Giant Branch (RGB) stars, have been performed by \citet{car06} and \citet{gra06} finding $\pm$ 0.01 and $\pm$ 0.04, respectively." + Carraroetal.(2006) also find about solar [|a /Fe]. while Grattonetal.(2006). find [O/Fe| depleted by a [actor of 2 with respect to the solar value.," \citet{car06} also find about solar $\alpha$ /Fe], while \citet{gra06} find [O/Fe] depleted by a factor of 2 with respect to the solar value." + The use of IR. spectroscopy offers an interesting alternative to optical spectroscopy. as it is less sensitive to the blanketing effects and more suitable to study cool and metal rich stars than the optical spectral range.," The use of IR spectroscopy offers an interesting alternative to optical spectroscopy, as it is less sensitive to the blanketing effects and more suitable to study cool and metal rich stars than the optical spectral range." + Our group has been undertaking a program using the NIRSPEC spectrograph (MeLean.19983). al Neck to obtain spectra of old metal rich stars in the bulge field (Rich&Origlia2005) and globular clusters (Origliaetal.2003:Origlia&Rich2004:Origlia.Valentietal.2005) with the aim of sliudvineg the composition and chemical evolution of the bulge aud globular clusters.," Our group has been undertaking a program using the NIRSPEC spectrograph \citep{ml98} + at Keck to obtain spectra of old metal rich stars in the bulge field \citep{ro05} and globular clusters \citep{ori03,ori04,or04,ori05} with the aim of studying the composition and chemical evolution of the bulge and globular clusters." + Precise chemical abundances of NGC 6791 are crucial to constrain better the age of this cluster. which deserves detailed investigations being one of the few examples in which we can study stars that formed very early in the evolution of the Galactic disk.," Precise chemical abundances of NGC 6791 are crucial to constrain better the age of this cluster, which deserves detailed investigations being one of the few examples in which we can study stars that formed very early in the evolution of the Galactic disk." +As underlined by,As underlined by +We caleulate the J7. on the Northern and Southern remispheres for a North pole at (6.6)=(57.10).,"We calculate the $J^3_\ell$ on the Northern and Southern hemispheres for a North pole at $(\ell,b)=(57,10)$." + In Fig., In Fig. + 5 we plot them against each other. ancl against the whole skv results.," \ref{N/S-J} we plot them against each other, and against the whole sky results." + We also plot the 1 sigma variance lines [or he whole sky. (the spread. is a little wider for a hall-sky)., We also plot the 1 sigma variance lines for the whole sky (the spread is a little wider for a half-sky). + We notice that the Northern results appear to take higher values than the Southern., We notice that the Northern results appear to take higher values than the Southern. + For 34 of the 49 points the vorthern result is higher. and when it is not the dillerence is small.," For 34 of the 49 points the Northern result is higher, and when it is not the difference is small." + We note that the individual V7 values do not reveal any significant departure from. Gaussianitv for either the vorthern or Southern results., We note that the individual $\chi^2$ values do not reveal any significant departure from Gaussianity for either the Northern or Southern results. + Lt is the dillerence between he hemispheres that is significant., It is the difference between the hemispheres that is significant. + We quantify this feature by measuring the sum of the cillerences: We calculate A for a North pole at (57.10). and at he ecliptic North pole position (96.30). and at a further 53 positions uniformly. distributed in the Northern galactic remisphere.," We quantify this feature by measuring the sum of the differences: We calculate $K$ for a North pole at (57,10), and at the ecliptic North pole position (96,30), and at a further 53 positions uniformly distributed in the Northern galactic hemisphere." + We infer the value of A for the opposite North »ole positions. in the Southern galactic hemisphere.," We infer the value of $K$ for the opposite North pole positions, in the Southern galactic hemisphere." + We then »erformi Gaussian Monte. Carlo simulations. as above. in order to evaluate the statistical significance of the observed WALAP asymmetry.," We then perform Gaussian Monte Carlo simulations, as above, in order to evaluate the statistical significance of the observed WMAP asymmetry." + For cach simulation we allow the North role to vary over the same positions and we record. the maximum [A] value seen in cach simulation., For each simulation we allow the North pole to vary over the same positions and we record the maximum $|K|$ value seen in each simulation. + Figure ο plots. at. the corresponding North pole »osition. the fraction of simulations with a lower |A[o value than the να) value from that North pole.," Figure \ref{K2} plots, at the corresponding North pole position, the fraction of simulations with a lower $|K|_{\it max}$ value than the $|K_{\it WMAP}|$ value from that North pole." + This raction is then displaved with the | sign of this Awaa result -l. and [1 indicates a North pole position returning higher asymmetry than all the simulations.," This fraction is then displayed with the $+/-$ sign of this $K_{\it WMAP}$ result -1, and +1 indicates a North pole position returning higher asymmetry than all the simulations." + We show the results for the ranges (=4100 (larger circles) and f=4—50 (smaller circles).," We show the results for the ranges $\ell=4-100$ (larger circles), and $\ell=4-50$ (smaller circles)." + In Table 2. we list the results for the three North pole positions that lind the maximum. value of A. for WMAT.," In Table \ref{Ktab} we list the results for the three North pole positions that find the maximum value of $K$, for WMAP." + We also show the fraction of simulations that return a lower LNuus value., We also show the fraction of simulations that return a lower $|K|_{\it max}$ value. + We observe significant asvmmetry at the ~99% level for both { ranges 4-100. ancl 4-50.," We observe significant asymmetry at the $\sim99\%$ level for both $\ell$ ranges 4-100, and 4-50." + The North pole positions that maximise this asymmetry ciller sliehtly between the { ranges. but both are consistent with the (57.10) position that Eviksen&al(2004a) [found maximises the angular power spectrum asvmnietry.," The North pole positions that maximise this asymmetry differ slightly between the $\ell$ ranges, but both are consistent with the (57,10) position that \cite{erik1} found maximises the angular power spectrum asymmetry." + Furthermore. the pattern on the sky displayed in Eig.," Furthermore, the pattern on the sky displayed in Fig." + 6 is very similar to that for the distribution of power. see Fig.," \ref{K2} is very similar to that for the distribution of power, see Fig." + 24 in Hansen.Banday&Córski(2004)., 24 in \cite{hbg}. +. Whether this anomaly is a sign of anisotropy. non-CGaussianitv or both is parth rhetorical.," Whether this anomaly is a sign of anisotropy, non-Gaussianity or both is partly rhetorical." + An anisotropic ensemble may always be converted into an isotropic one by randomising the preferred axis., An anisotropic ensemble may always be converted into an isotropic one by randomising the preferred axis. + From the point of view of such an ensemble. what appears in a single skv as a sign of anisotropy is simply non-Gaussianitv (see. Ferreira&Magueijo(1997). [or details).," From the point of view of such an ensemble, what appears in a single sky as a sign of anisotropy is simply non-Gaussianity (see \cite{aniso} for details)." + As previously mentioned. for a Gaussian random: field the normalised bispectrum and the power spectrum are independent quantities.," As previously mentioned, for a Gaussian random field the normalised bispectrum and the power spectrum are independent quantities." + In the analysis in this section we were particularly careful to avoid selection effects (we let. Gaussian simulations choose their maximal asvnimectry axis)., In the analysis in this section we were particularly careful to avoid selection effects (we let Gaussian simulations choose their maximal asymmetry axis). + Should one insist that the anomaly found is a Luke. however. we should stress that it would be an independent Juke with respect to the anomaly found. in the power spectrum (Iriksen&al2004a).," Should one insist that the anomaly found is a fluke, however, we should stress that it would be an independent fluke with respect to the anomaly found in the power spectrum \citep{erik1}." +. This makes it even more unlikely., This makes it even more unlikely. + We shall return to this matter in Section ??.., We shall return to this matter in Section \ref{select}. + We caleulate the single-£ normalised bispectrum. £7. on the Northern and. Southern hemispheres for a North pole at ((.b) =(57.10).," We calculate the $\ell$ normalised bispectrum, $I^3_\ell$, on the Northern and Southern hemispheres for a North pole at $(\ell,b)=$ (57,10)." + In Fie., In Fig. + T we plot them against each other. and against the whole sky results.," \ref{N/S-I} we plot them against each other, and against the whole sky results." + We also plot the 1 sigma error bars for the whole sky (the spread. is very. similar for a hallsky)., We also plot the 1 sigma error bars for the whole sky (the spread is very similar for a half-sky). + The Northern results appear to Luctuate more, The Northern results appear to fluctuate more +"Tjese new rates. and others from ΕΠΙ uot explicitly liseussed above, were incorporated im the uew solar Mocel calculations.","These new rates, and others from SFII not explicitly discussed above, were incorporated in the new solar model calculations." + A] iiodoels in this work are required to reproduce the reselt-¢av (7.1.57 Cyr) solar luminosity. L.BALLS$, is also slightly better in the new models." + The changes are. however. small compared to mo uncertainties. so the new models are effectively fulv consistent with previous ones.," The changes are, however, small compared to model uncertainties, so the new models are effectively fully consistent with previous ones." + Given the similarities between the new models axd hose of Serenellietal.(2009)... the reader is referred ο that source for a more detailed discussion.," Given the similarities between the new models and those of \citet{ssm:09}, the reader is referred to that source for a more detailed discussion." + Iere. we cus on manifestatious of the solar abundance proble11. he primary motivation for the studies presented her," Here, we focus on manifestations of the solar abundance problem, the primary motivation for the studies presented here." + There ave three key indicators of the abundauce proble he depth of the couvective cuvelope. the surface heli abundance. aud the sound speed profile.," There are three key indicators of the abundance problem, the depth of the convective envelope, the surface helium abundance, and the sound speed profile." +" The ACSSY) nodel produces a convective zone that is too shall""v and aos rface helm abuudauce that is too low (see Table 1)).", The AGSS09 model produces a convective zone that is too shallow and a surface helium abundance that is too low (see Table \ref{tab:ssm}) ). + The cutries in the table should be compared to the hejoscisuuc values. Rey=0.7154FE0.001 (Basa&Autia1997) and 0.2[85+0.0035 (Basu&Antia2001).. respectively.," The entries in the table should be compared to the helioseismic values, $\rcz= 0.713\pm 0.001$ \citep{basu:1997} and $0.2485 \pm +0.0035$ \citep {basu:2004}, , respectively." + If liclioscisuic uncertainties are combined with SSAL uncertaimties (Baheallctal.2006... equations 32 and 35). then ACGSSO9 model predictions differ from solar values at Jo for oth. Rey aud Vy.," If helioseismic uncertainties are combined with SSM uncertainties \citealp {montecarlo:2006}, equations 32 and 35), then AGSS09 model predictions differ from solar values at $3\sigma$ for both $\rcz$ and $\ys$." + In contrast. CS98 model precdictious are in excellent aerecineut witli he parameters extractec from Lelioscismoloey.," In contrast, GS98 model predictions are in excellent agreement with the parameters extracted from helioseismology." +" Altloteh nore difficult to «uantifv. the difference between he AGSSO9 and solar ""OIud speed profiles. determined your (efc). is aout four times larger than tlat obtained with the higher metalliciv GS98 model."," Although more difficult to quantify, the difference between the AGSS09 and solar sound speed profiles, determined from $\left< \delta c/c \right>$, is about four times larger than that obtained with the higher metallicity GS98 model." + These specific discrepancies for the ACSSOO abuudauce set are he quantitative manifestations othe solar abunda1ος xoblem metioned in the introduction., These specific discrepancies for the AGSS09 abundance set are the quantitative manifestations of the solar abundance problem mentioned in the introduction. +" The new SFIDI nuclear reaction rates do alter he predicted solar neutrino fhNOS, warticularly he Huxes associated with the ppII cin and €NLevele. iiechauisuis for solar Ἠνοτεech burning that are relatively unimportant cherecticalve"," The new SFII nuclear reaction rates do alter the predicted solar neutrino fluxes, particularly the fluxes associated with the ppIII chain and CN-cycle, mechanisms for solar hydrogen burning that are relatively unimportant energetically." + Wihn respect to he rates used in Sereuellietal. (2009).. the lost inportaut SEII changes are the Πιοοαπο 1n S41(0). the decrease in S35(0). aud the ΠιοCase dl 544400).," With respect to the rates used in \citet{ssm:09}, , the most important SFII changes are the increase in ${\rm S_{11}(0)}$, the decrease in ${\rm S_{33}(0)}$, and the increase in ${\rm S_{1\, 14}(0)}$." + Otos changes affecting predicted ueutrino fntxes Include he reduced proton capture cross section ¢m DN aud he increase in the pep/ op rate ratio., Other changes affecting predicted neutrino fluxes include the reduced proton capture cross section on $^{15}$ N and the increase in the pep/ pp rate ratio. + Mkxlel predictious or the neutriuo fluxes and associated 1necrtaintics are xeseuted in Table 2 for both the ACSSO00 and CS98 conrpositious., Model predictions for the neutrino fluxes and associated uncertainties are presented in Table \ref{tab:neutrinos} for both the AGSS09 and GS98 compositions. + The nios sienificant changes are a decrease in the prediced *B flux primarily because of he ierease m S44 21x hie increase in the LN flux due o the larger ceutral aburdance of €. The increase in € is a cousequence of the lower SFIT value for ΝΟ 10. a reaction that competes with the CN I cvele reaction Np.) C aud allows nass to flowout of the CN Devele iuto CN IL.," The most significant changes are a decrease in the predicted $^8$ B flux primarily because of the increase in ${\rm + S_{11}}$ and the increase in the $^{13}$ N flux due to the larger central abundance of C. The increase in C is a consequence of the lower SFII value for $^{15}$ $\gamma)^{16}$ O, a reaction that competes with the CN I cycle reaction $^{15}$ $\alpha)^{12}$ C and allows mass to flowout of the CNI cycle into CN II." + We discuss t1e parainctricdependence of the fixes on nuclearcross sections and other SSAL iuput xuwanmeters elsewhere. including some of the resulting jeutrino tests that might be made of solarcomposition aud S-factors (Serenellietal.2011 |.," We discuss the parametricdependence of the fluxes on nuclearcross sections and other SSM input parameters elsewhere, including some of the resulting neutrino tests that might be made of solarcomposition and S-factors \citep{PGSH}. ." +eiven in the literature.,given in the literature. + We made the measurements on the average of spectra summed to one hour long exposure times. but we plotted the daily average values of the EWs on final ciagrams. shown in Fig. 6..," We made the measurements on the average of spectra summed to one hour long exposure times, but we plotted the daily average values of the EWs on final diagrams, shown in Fig. \ref{abs_EW}," + to better see the trends., to better see the trends. + For the Ho line we found that EW changes from 2.50 to 2.85 oover an orbital phase (see Fig. 6))., For the $\alpha$ line we found that EW changes from 2.50 to 2.85 over an orbital phase (see Fig. \ref{abs_EW}) ). + The average value of 270A aagrees very well. within the uncertainties. with the resultof €05 (2.8 d: 0.1 AJ). except they found that the value was stable during the orbit of the binary (which could possibly be explained by the relatively low resolution of their spectra).," The average value of 2.70 agrees very well, within the uncertainties, with the resultof C05 (2.8 $\pm$ 0.1 ), except they found that the value was stable during the orbit of the binary (which could possibly be explained by the relatively low resolution of their spectra)." + Quir result is also consistent with the EW values measured by others over the last ten vears (Bosch-Ramonetal.2007).., Our result is also consistent with the EW values measured by others over the last ten years \citep{Bosch-Ramon2007}. . + Using the method. of Pulsetal.(1996). we estimated the mass loss rate [rom the EW of the Ho line., Using the method of \citet{Puls1996} we estimated the mass loss rate from the EW of the $\alpha$ line. + To do these calculations we adopted. the following parameters: Ro = Q5. Re and für = 39.000 + 1000 IX for the O-type star (C05): a terminal wine velocity of V4 = 2440 + 190 kms * with a wind velocity law exponent of 7 = 0.8 (MOA).," To do these calculations we adopted the following parameters: $R_{\rm O}$ = $^{+0.7}_{-0.6}$, $_{\odot}$ and $T_{\rm eff}$ = 39,000 $\pm$ 1000 K for the O-type star (C05); a terminal wind velocity of $V_{\infty}$ = 2440 $\pm$ 190 km $^{-1}$ with a wind velocity law exponent of $\beta$ = 0.8 (M04)." +" We found that the mass-loss rate of the stellar companion is around 3.7 10* M. vr.| from the strongest absorption. which corresponds to the lower limit. and 4.8 "" | for the upper limit."," We found that the mass-loss rate of the stellar companion is around 3.7 $\times$ $^{-7}$ $_{\odot}$ $^{-1}$ from the strongest absorption, which corresponds to the lower limit, and 4.8 $\times$ $^{-7}$ $_{\odot}$ $^{-1}$ for the upper limit." +" These values are consistent with he mass loss rates of z3.7610 M. 1 for the low state (strong absorption) and zz7.5.10""AL. + for the uigh state obtained by 05 (see also MOA).", These values are consistent with the mass loss rates of $\approx 3.7 \times 10^{-7}$ $_{\odot}$ $^{-1}$ for the low state (strong absorption) and $\approx 7.5 \times 10^{-7}$ $_{\odot}$ $^{-1}$ for the high state obtained by C05 (see also M04). + We found two lines (LL? and. He A5SS75) showing significant changes curing the orbit., We found two lines $\beta$ and He $\lambda$ 5875) showing significant changes during the orbit. + Phe lowest absorption or the LE2 line occurs around y~ 0.75. and at ye~ 0.65 for he He line refabsgW 3). efosetothecepectedphases ~ O.7 when the compact object is between us and. the stellar. companion (inferior conjunction).," The lowest absorption for the $\beta$ line occurs around $\varphi\sim$ 0.75, and at $\varphi\sim$ 0.65 for the He line \\ref{abs_EW}) ), close to the expected phase $\varphi\sim$ 0.7 when the compact object is between us and the stellar companion (inferior conjunction)." + Wo carricd out a simple correlation analysis between these EW changes and theAJOST ligh curvesize). and found the correlation coellicients. r. to be 0.52. 0.70anc 0.52. for the Ho. Ll? and He A5875 lines. respectively.," We carried out a simple correlation analysis between these EW changes and the light curve, and found the correlation coefficients, $r$, to be 0.52, 0.70and 0.52 for the $\alpha$, $\beta$ and He $\lambda$ 5875 lines, respectively." + This resul sugeests the possibility of real mocdulations over the orbita period. but further studies are necessary.," This result suggests the possibility of real modulations over the orbital period, but further studies are necessary." + Note that our. data only show smooth orbita modulation in the tla. Lh? and Le absorption lines ane in the emission components of the LL Balmer lines.," Note that our data only show smooth orbital modulation in the $\alpha$, ${\beta}$ and He absorption lines and in the emission components of the H Balmer lines." + The lines do not show evidence of dense clump-like condensates in the stellar wind., The lines do not show evidence of dense clump-like condensates in the stellar wind. + The two main findings from our work are that the orbit has a significantly. lower cecentricity than previously thought and that the The eccentricity ο=0.4 obtained by MOA put LS 5039 as the system with the most eccentric orbit among X-ray. binaries with an © donor star., The two main findings from our work are that the orbit has a significantly lower eccentricity than previously thought and that the The eccentricity $e\approx 0.4$ obtained by M04 put LS 5039 as the system with the most eccentric orbit among X-ray binaries with an O donor star. + Although the value was later shown to be lower (cz 0.35. C05 and. A09). the eccentricity would.still be," Although the value was later shown to be lower $e\approx 0.35$ , C05 and A09), the eccentricity wouldstill be" +For the correlations between BH mass and central stellar velocity dispersion. one obtains 2.0 x10*A£. for NGC 5077 adopting the form by Tremaineetal.(2002) (see 13.. upper panel). a factor 3.4 lower than our estimate.,"For the correlations between BH mass and central stellar velocity dispersion, one obtains 2.0 $\times 10^8 M_{\odot}$ for NGC 5077 adopting the form by \citet{tremaine02} (see \ref{Mbh_sigma}, upper panel), a factor 3.4 lower than our estimate." + We also compared our measurement with the expected value from the gy—crelation derived by Ferrarese&Ford(2005)... Mpy = 5.4 x109M. in full agreement with our estimate (see 13.. bottom panel).," We also compared our measurement with the expected value from the $M_{\rm BH}-\sigma$relation derived by \citet{ferrarese05}, $M_{\rm BH}$ = 5.4 $\times 10^{8} M_{\odot}$, in full agreement with our estimate (see \ref{Mbh_sigma}, bottom panel)." + Conversely. the BH mass expected for NGC 5077 on the basis of its correlation with the Sérrsic concentration index (fromTrujilloetal.2004) Is a factor 5-6 lower than our estimate when using the linear Mpg- Sérrsic index correlation (Graham&Driver2007) anda factor 3 lower when using the quadratic form they proposed.," Conversely, the BH mass expected for NGC 5077 on the basis of its correlation with the Sérrsic concentration index \citep[from][]{trujillo04} is a factor $\sim$ 6 lower than our estimate when using the linear $M_{\rm BH} -$ Sérrsic index correlation \citep{graham07} and a factor 3 lower when using the quadratic form they proposed." + Marconi&Hunt(2003) noticed that Mgy is separately correlated with both σ and Ry., \citet{marconi03} noticed that $M_{\rm BH}$ is separately correlated with both $\sigma$ and $R_{\rm e}$. + This is shown by the weak correlation between the residuals of the Mpg—c correlation with Ας reproduced here in Fig., This is shown by the weak correlation between the residuals of the $M_{\rm BH}-\sigma$ correlation with $R_{\rm e}$ reproduced here in Fig. + 14. for the Tremaine at al., \ref{residuals} for the Tremaine at al. +form?.. Our determinations of the black hole mass in NGC 5077 and in NGC 3998 (Paper I) support this idea., Our determinations of the black hole mass in NGC 5077 and in NGC 3998 (Paper I) support this idea. +" Unlike NGC 3998. which has one of the lowest values of R, among galaxies with measured Mgg (0.85 kpe) and shows a negative residual from the Mgy-c correlation. NGC 5077 has an intermediate R, value (3.6 kpe) and show a small but positive residual (see Fig. 14))."," Unlike NGC 3998, which has one of the lowest values of $R_{\rm e}$ among galaxies with measured $M_{\rm BH}$ (0.85 kpc) and shows a negative residual from the $M_{\rm BH}-\sigma$ correlation, NGC 5077 has an intermediate $R_{\rm e}$ value (3.6 kpc) and show a small but positive residual (see Fig. \ref{residuals}) )." + In the same sense. but with a larger positive residual. there is the result found by Capettietal.(2005) for NGC 532535. ἃ galaxy with quite a large effective radius (9.7 kpe).," In the same sense, but with a larger positive residual, there is the result found by \citet{Capetti05} for NGC 5252, a galaxy with quite a large effective radius (9.7 kpc)." +" This indicates the presence of a black hole’s ""fundamental plane? in the sense that a combination of at least « and R. drives the correlations between μη and the bulge properties.", This indicates the presence of a black hole's “fundamental plane” in the sense that a combination of at least $\sigma$ and $R_{\rm e}$ drives the correlations between $M_{\rm BH}$ and the bulge properties. + The physical implications of these results will be discussed in a forthcoming paper. when new direct BH mass measurements are presented. in order to base our discussion on a statistically more significative sample of M3y determinations.," The physical implications of these results will be discussed in a forthcoming paper, when new direct BH mass measurements are presented, in order to base our discussion on a statistically more significative sample of $M_{\rm BH}$ determinations." +" We have presented results from à gas kinematics study in the nucleus of three nearby LINER galaxies: IC 989, NGC 5077. and NGC 6500 using archival HST/STIS spectra."," We have presented results from a gas kinematics study in the nucleus of three nearby LINER galaxies: IC 989, NGC 5077, and NGC 6500 using archival HST/STIS spectra." + Only in the case of NGC 5077 does the nuclear velocity curves appear to be associated with gas in regular rotation., Only in the case of NGC 5077 does the nuclear velocity curves appear to be associated with gas in regular rotation. + For IC 989 our results indicate an inner counter-rotating gas system. while for NGC 6500 the complex trend in the velocity curves suggests à nuclear expanding bubble.," For IC 989 our results indicate an inner counter-rotating gas system, while for NGC 6500 the complex trend in the velocity curves suggests a nuclear expanding bubble." +" We used our modeling code to fit the observed [N II],16583 surface brightness distribution and velocity curve of NGC 5077.", We used our modeling code to fit the observed [N $\lambda$ 6583 surface brightness distribution and velocity curve of NGC 5077. + The dynamics of the rotating gas can be accurately reproduced by motions in a thin disk when a compact dark mass of Mp3y=6.813x10°M. is added to the stellar mass component.," The dynamics of the rotating gas can be accurately reproduced by motions in a thin disk when a compact dark mass of $M_{\rm BH} = +6.8_{-2.8}^{+4.3}\times 10^8 M_{\odot}$ is added to the stellar mass component." + Furthermore. the black hole in NGC 5077 has a sphere of influence radius. Ryn=σΜιμμ/c0 of ~ 62 pe (= 0734). well-resolved at the HST resolution (2Run/Rf5 6.8).," Furthermore, the black hole in NGC 5077 has a sphere of influence radius, $R_{\rm sph} = GM_{\rm BH}/\sigma_{\rm star}^{2}$ , of $\sim$ 62 pc $\simeq$ $0\farcs34$ ), well-resolved at the HST resolution $R_{\rm sph}/R_{\rm res} \simeq$ 6.8)." + For what concerns the connections of this BH mass estimate with the properties of the host galaxy. the Mgy value," For what concerns the connections of this BH mass estimate with the properties of the host galaxy, the $M_{\rm BH}$ value" +the solid and dashed black lines in Fig. 2..,the solid and dashed black lines in Fig. \ref{fig:xi}. + This means that the large-scale shape of the correlation function is essentially driven by the underlying mass distribution in the assumed cosmological model and not by the details of the semi-analytic recipe adopted to generate galaxies., This means that the large-scale shape of the correlation function is essentially driven by the underlying mass distribution in the assumed cosmological model and not by the details of the semi-analytic recipe adopted to generate galaxies. +" A different recipe would not affect, therefore, the results obtained here, unless we postulate the existence of dramatically non-local galaxy formation processes (e.g.?).."," A different recipe would not affect, therefore, the results obtained here, unless we postulate the existence of dramatically non-local galaxy formation processes \citep[e.g.][]{narayanan00}." + We also note from Fig., We also note from Fig. + 2 that the dependence of on the PDF threshold is essentially on large scales., \ref{fig:xi} that the dependence of on the PDF threshold is essentially on large scales. + Below ~1h~' there is no significant change when denser and denser environments are excluded., Below $\sim 1$ there is no significant change when denser and denser environments are excluded. +" In their analysis of the SDSS ? consider sub-samples defined as extrema of the density distribution, i.e. using galaxies lying on the tails of the distribution on both sides."," In their analysis of the SDSS \citet{abbas07} consider sub-samples defined as extrema of the density distribution, i.e. using galaxies lying on the tails of the distribution on both sides." +" With this selection, they find a change in for different environments also on small scales."," With this selection, they find a change in for different environments also on small scales." +" It can be shown simply using the conservation of galaxy pairs (see?) that the two results are in fact consistent with each other (Ravi Sheth, private communication)."," It can be shown simply using the conservation of galaxy pairs \citep[see Eq. 1 of ][]{abbas07} that the two results are in fact consistent with each other (Ravi Sheth, private communication)." +" These results highlight the importance in redshift surveys of an accurate reconstruction of the density field, to evidence possible peculiarities in the overall PDF as sampled by that specific catalogue."," These results highlight the importance in redshift surveys of an accurate reconstruction of the density field, to evidence possible peculiarities in the overall PDF as sampled by that specific catalogue." +" Further strengthening the results obtained by ? at zc0, we have shown that an anomalous density distribution function can significantly bias the recovered two-point correlation function, making it difficult to draw general conclusions from its shape."," Further strengthening the results obtained by \citet{abbas07} at $z\simeq 0$, we have shown that an anomalous density distribution function can significantly bias the recovered two-point correlation function, making it difficult to draw general conclusions from its shape." + This result provides another example of the intrinsic difficulty existing when comparing observations of the galaxy distribution to theoretical predictions., This result provides another example of the intrinsic difficulty existing when comparing observations of the galaxy distribution to theoretical predictions. + The theory provides us with fairly accurate forecasts for the distribution of the dark matter and for that of the halos within which we believe galaxies form (e.g.??)..," The theory provides us with fairly accurate forecasts for the distribution of the dark matter and for that of the halos within which we believe galaxies form \citep[e.g.][]{mo96,sheth99}." +" However, translating galaxy clustering measurements into constraints for the halo clustering involves understanding how the selected galaxies populate halos with different mass."," However, translating galaxy clustering measurements into constraints for the halo clustering involves understanding how the selected galaxies populate halos with different mass." +" The result presented here show how a sample particularly rich in dense structures favours higher-mass halos, which in turn are more clustered, thus biasing the observed correlation function as a function of scale."," The result presented here show how a sample particularly rich in dense structures favours higher-mass halos, which in turn are more clustered, thus biasing the observed correlation function as a function of scale." + A more detailed analysis of the environmental dependence of galaxy clustering in the zCOSMOS-Bright sample and related HOD modelling will be presented in a future paper., A more detailed analysis of the environmental dependence of galaxy clustering in the zCOSMOS-Bright sample and related HOD modelling will be presented in a future paper. + We thank Ravi Sheth for helpful comments on the manuscript., We thank Ravi Sheth for helpful comments on the manuscript. + Financial support from INAF and ASI through grants PRIN-INAF-2007 and ASI/COFIS/WP3110 1/026/07/0 is gratefully acknowledged., Financial support from INAF and ASI through grants PRIN-INAF--2007 and ASI/COFIS/WP3110 I/026/07/0 is gratefully acknowledged. +" LG thanks D. Sanders and the University of Hawaii, for hospitality at the Institute for Astronomy, where this work was initiated."," LG thanks D. Sanders and the University of Hawaii, for hospitality at the Institute for Astronomy, where this work was initiated." + This work is based on observations undertaken at the European Southern Observatory (ESO) Very Large Telescope (VLT) under Large Program 175.A-0839., This work is based on observations undertaken at the European Southern Observatory (ESO) Very Large Telescope (VLT) under Large Program 175.A-0839. +" Also based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, operated by the Association of Universities for Research in Astronomy, Inc. (AURA), under NASA contract NAS 5Y26555, with the Subaru Telescope, operated by the National Astronomical Observatory of Japan, with the telescopes of the National Optical Astronomy Observatory, operated by the Association of Universities for Research in Astronomy, Inc. (AURA), under cooperative agreement with the National Science Foundation, and with the Canada-France- Hawaii Telescope, operated by the National Research Council of Canada, the Centre National de la Recherche Scientifique de France, and the University of Hawaii."," Also based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, operated by the Association of Universities for Research in Astronomy, Inc. (AURA), under NASA contract NAS 5Y26555, with the Subaru Telescope, operated by the National Astronomical Observatory of Japan, with the telescopes of the National Optical Astronomy Observatory, operated by the Association of Universities for Research in Astronomy, Inc. (AURA), under cooperative agreement with the National Science Foundation, and with the Canada-France- Hawaii Telescope, operated by the National Research Council of Canada, the Centre National de la Recherche Scientifique de France, and the University of Hawaii." +instantaneous (ICGsr) burst of star formation.,instantaneous $\sim1$ Gyr) burst of star formation. + Acopting the νι burst. model of Bruzual Charlot (1993) (updated. 19905 models). we show in big.," Adopting the 1Gyr burst model of Bruzual Charlot \shortcite{Bruzual1993} (updated 1995 models), we show in Fig." + 2 the spectral evolution as a function. of age since the initial starburst., \ref{specevol} the spectral evolution as a function of age since the initial starburst. +" Also shown beside cach age is the approximate redshift for a formation epoch of 2,=5.", Also shown beside each age is the approximate redshift for a formation epoch of $z_{f}=5$. + Initially. the 40004 breakregionisdiluledbystarburstactivilg.," Initially, the $\,$ break region is diluted by starburst activity." + Aflera f ewgigagears.thestellarpopulalionecolvespassiccly.maintaininganalmostuniform*t ," After a few gigayears, the stellar population evolves passively, maintaining an almost uniform SED shape (see Bruzual Charlot 1993)." +l.," For formation epochs $z\simgt5$, it is apparent that the generic elliptical SED shape of Fig." + 1, \ref{galspec} is easily established by $z\sim1$ . +n general. the total flux at a given wavelength observed in a Parkes source. f(A). can be modelled as the sum of light contributed by the central quasar or AGN. f(A). and any underlying host galaxy f(A) (eg.," In general, the total flux at a given wavelength observed in a Parkes source, $f_{T}(\lambda)$, can be modelled as the sum of light contributed by the central quasar or AGN, $f_{q}(\lambda)$, and any underlying host galaxy $f_{g}(\lambda)$ (eg." + Pig. 1))., Fig. \ref{galspec}) ). + We write the relationship between these quantities as: where e is a scaling factor giving an arbitrary measureof the amount of galaxy light we wish to determine for a particular SOULCE., We write the relationship between these quantities as: where $c$ is a scaling factor giving an arbitrary measureof the amount of galaxy light we wish to determine for a particular source. + Given fr(A) for a particular source ancl an arbitrary galaxy spectrum fy(A) (Fig. 1)).," Given $f_{T}(\lambda)$ for a particular source and an arbitrary galaxy spectrum $f_{g}(\lambda)$ (Fig. \ref{galspec}) )," + our aim is to determine the value of e such that f(A)looks something like a quasar spectrum., our aim is to determine the value of $c$ such that $f_{q}(\lambda)$looks something like a quasar spectrum. +" From our knowledge of quasar optical spectra. an obvious choice is to require that f(A) be. ""smooth? and contain no breaks."," From our knowledge of quasar optical spectra, an obvious choice is to require that $f_{q}(\lambda)$ be “smooth” and contain no breaks." +" Thus. the basis of this algorithm involves subtracting an arbitrary amount of galaxy. [lux. cfA) from f(A) such that the resulting spectrum |f,(À) appears smooth (see below)."," Thus, the basis of this algorithm involves subtracting an arbitrary amount of galaxy flux, $c\,f_{g}(\lambda)$ from $f_{T}(\lambda)$ such that the resulting spectrum $f_{q}(\lambda)$ appears smooth (see below)." + When this is achieved. the fraction of total light at a given wavelength contributed by the host galaxy can be estimated by normalising: In order to implement the above algorithm. we need to define an acceptable form for the shape of the quasar spectrum fyfA).," When this is achieved, the fraction of total light at a given wavelength contributed by the host galaxy can be estimated by normalising: In order to implement the above algorithm, we need to define an acceptable form for the shape of the quasar spectrum $f_{q}(\lambda)$." + Our only requirement is that this spectrum be smooth and hence our choice is somewhat arbitrary., Our only requirement is that this spectrum be smooth and hence our choice is somewhat arbitrary. +" We choose f(A) to be a power-law. paramoeterised by: where the slope à ids determined. between. (two. fixed wavelengths. À,,;, and Ai; (see below): Using Eqn. 1.."," We choose $f_{q}(\lambda)$ to be a power-law, parameterised by: where the slope $\alpha$ is determined between two fixed wavelengths, $\lambda_{min}$ and $\lambda_{max}$ (see below): Using Eqn. \ref{fq}," + the Ηχος defined in Eqn., the fluxes defined in Eqn. +" 4. can be written: For a discussion on how the wavelengths À,,;, and Aner are chosen and the fluxes in Eqn.", \ref{alpha2} can be written: For a discussion on how the wavelengths $\lambda_{min}$ and $\lambda_{max}$ are chosen and the fluxes in Eqn. + 5. measured. see section 3.1..," \ref{fqminmax} measured, see section \ref{apal}." + Phe validity of our assumption of a single for f(A) is discussed in section 6.., The validity of our assumption of a single power-law for $f_{q}(\lambda)$ is discussed in section \ref{diseven}. +" To apply this algorithm in a self. consistent way to each of our optical spectra. we need to define a Ligure of merit indicating the point at which the maximum amount of galaxy spectrum has been subtractecl and a ""wmooth quasar spectrum fie.α power-law) is achieved."," To apply this algorithm in a self consistent way to each of our optical spectra, we need to define a figure of merit indicating the point at which the maximum amount of galaxy spectrum has been subtracted and a “smooth” quasar spectrum (ie.a power-law) is achieved." + Let us. first. consider the rest. frame optical spectrum of a source. f(A). suspected. of containing a 4000A break.," Let us first consider the rest frame optical spectrum of a source, $f_{T}(\lambda)$, suspected of containing a $\,$ break." + Ehisisillustratedinl ig. 3., This is illustrated in Fig. \ref{schematic}. + Purlhermore.lebusconsidera7smoo law(ie," Furthermore, let us consider a “smooth” power-law (ie." +" theunderlyinggquasarspecdriml(A)) between two wavelengths À,,;, and Aya, on either side. of the 4000A breakregionasshouwn"," the underlying quasar spectrum $f_{q}(\lambda)$ ) between two wavelengths $\lambda_{min}$ and $\lambda_{max}$ on either side of the $\,$ break region as shown." + M edefincourfigureof merilascrepresentinglhe , We define our figure of merit as representing the area $A$ of the shaded region in Fig. \ref{schematic}. +where f4(A) is given by Fie. 12) ," As certain amounts of galaxy spectrum, $c\,f_{g}(\lambda)$ (where $f_{g}(\lambda)$ is given by Fig. \ref{galspec}) )" +are gradually. subtracted from f(A). 2 will decrease and becomes a minimum when 10 break disappears.," are gradually subtracted from $f_{T}(\lambda)$, $A$ will decrease and becomes a minimum when the break disappears." + We can thus determine the value of c when this occurs. allowing us to estimate the fractional galaxy contribution from Eqn. 2..," We can thus determine the value of $c$ when this occurs, allowing us to estimate the fractional galaxy contribution from Eqn. \ref{galfrac}." + For à given amount of subtracted galaxy Dux. efy(A) we can write ;lin terms of e from Eqn.," For a given amount of subtracted galaxy flux, $c\,f_{g}(\lambda)$, we can write $A$ in terms of $c$ from Eqn." + |. as follows: With JuifA) defined bv Eqns. 3.. ," \ref{fq} as follows: With $f_{q}(\lambda)$ defined by Eqns. \ref{pl2}, ," +4. and. 5.. ch can be written: where ale.fr.fy) is defined by Eqns.," \ref{alpha2} and \ref{fqminmax}, , $A$ can be written: where $\alpha(c,f_{T},f_{g})$ is defined by Eqns." + 4 and. 5.., \ref{alpha2} and \ref{fqminmax}. + Thus. one would only have to minimise zl with respect to c in order to determine the maximal ealactic contribution.," Thus, one would only have to minimise $A$ with respect to $c$ in order to determine the maximal galactic contribution." + We introduce however. an additional factor in Eqn.," We introduce however, an additional factor in Eqn." +7 whose purpose will be to make best use of the available data and optimise our algorithm.,\ref{Atwo} whose purpose will be to make best use of the available data and optimise our algorithm. +offer a large range of masses with similar ages and distances.,offer a large range of masses with similar ages and distances. + Muench.Lada.&Lada(2000) used the Baralleetal.(1993). models to perform a similar analvsis. to masses as low as 0.02 AL.. and probed a complementary age range (< LO Myr) to the current work.," \citet{mll} used the \citet{bcah98} models to perform a similar analysis, to masses as low as 0.02 $_{\odot}$, and probed a complementary age range $<$ 10 Myr) to the current work." + The outline of this paper is as follows: 82 provides a summary of the input evolutionary models and bolometric corrections for our analvsis: 83 describes (he main features of the predicted luminosity functions., The outline of this paper is as follows: $\S$ 2 provides a summary of the input evolutionary models and bolometric corrections for our analysis; $\S$ 3 describes the main features of the predicted luminosity functions. + In 84. we compare (those results against observations of several voung clusters. and consider the potential offered by future observations.," In $\S$ 4, we compare those results against observations of several young clusters, and consider the potential offered by future observations." + Most current stucdies of very low-mass dwarls are based on one of two sets of theoretical models., Most current studies of very low-mass dwarfs are based on one of two sets of theoretical models. + Baralleetal.(1998) combine interior models with detailed model atmosphere calculations. but cover only a limited range of mass and effective temperature (essentially corresponding to a lower spectral tvpe limit of L): Chabrier.Baralle.(2000) update these models to cover a lower range of masses. 0.1 M. to 0.01 M. (spectral tvpes M. L. and T).," \citet{bcah98} combine interior models with detailed model atmosphere calculations, but cover only a limited range of mass and effective temperature (essentially corresponding to a lower spectral type limit of L); \citet{cbah00} update these models to cover a lower range of masses, 0.1 $_{\odot}$ to 0.01 $_{\odot}$ (spectral types M, L, and T)." +" In contrast. Burrowsetal.(2001) present a more extensive series of models. spanning masses between 1, and 0.15AL. (1AZ.=1047 Mj) and ages in the range 10* to LOM vears. that do not include model atimospheres. but predict bolometric Iuminosities and effective temperatures as a function of time."," In contrast, \citet{bur} present a more extensive series of models, spanning masses between $1~M_J$ and $0.15~M_{\odot}$ $1~M_{\odot} = 1047~M_J$ ) and ages in the range $10^6$ to $10^{10}$ years, that do not include model atmospheres, but predict bolometric luminosities and effective temperatures as a function of time." + Since our aim is (o include coverage of very low-nmass stars through low-temperature brown dwarls. we base our analvsis on the moclels.," Since our aim is to include coverage of very low-mass stars through low-temperature brown dwarfs, we base our analysis on the \citet{bur} models." + Figure 1 shows the (me variation in elfective temperature predicted by the Burrows models., Figure 1 shows the time variation in effective temperature predicted by the Burrows models. + These models provide a series of reference points for specilic masses αἱ particular ages., These models provide a series of reference points for specific masses at particular ages. + To achieve the numerical resolution needed [or our study. we interpolate among the predicted values of bolometric luminosity and effective temperature using a bicubic spline interpolation as a hunction of mass and age.," To achieve the numerical resolution needed for our study, we interpolate among the predicted values of bolometric luminosity and effective temperature using a bicubic spline interpolation as a function of mass and age." + This interpolation steps over fine increments in both mass (Am=ἐκ] 1ÀU.) and age (Nlog(Age(Gyvr))=8.602x10. 1)., This interpolation steps over fine increments in both mass ${\Delta}m = 1{\times}10^{-4}M_{\odot}$ ) and age ${\rm \Delta\log (Age(Gyr)) = 8.602{\times}10^{-4}}$ ). + We do nol altempt to model clusters with ages of less than 23 Myr. since the model calculations are unreliable at these ages (Burrowsetal.2001).," We do not attempt to model clusters with ages of less than 2–3 Myr, since the model calculations are unreliable at these ages \citep{bur}." +. Qur goal is to compare (he model precictions against observations of near- ancl (14H NM) luminosity finetions of voung stellar clusters., Our goal is to compare the model predictions against observations of near- and mid-infrared $IJHKM$ ) luminosity functions of young stellar clusters. + As a first step. bolometric lhuninositv funcüons are generated.," As a first step, bolometric luminosity functions are generated." + This process requires assumptions concerning both the IMF and the distribution in age., This process requires assumptions concerning both the IMF and the distribution in age. + We parameterize (he IME. as a two-seement power-law. (ay.ao).," We parameterize the IMF as a two-segment power-law, $\alpha_1, \alpha_2$ )." + The higher-mass segment includes the range 0.08AZ. O.15 ALL. with ay=1.05 ," The higher-mass segment includes the range $0.08~M_{\odot}$ $0.15~M_{\odot}$ , with $\alpha_1 = 1.05$ " +Hy and isotopologues.,$_3^+$ and isotopologues. + The ratio increases with temperature in this temperature range. as already explained in detail by?.," The ratio increases with temperature in this temperature range, as already explained in detail by." +. The three thick lines are the ratios as calculated with RADEX from the observed line intensities (as a function of temperature and for different densitiesand o/p-H» ratios. as already shown in Fig. 6)).," The three thick lines are the ratios as calculated with RADEX from the observed line intensities (as a function of temperature and for different densitiesand $_2$ ratios, as already shown in Fig. \ref{ratio}) )." + We discuss in the following section the comparison between the chemical model and the observations., We discuss in the following section the comparison between the chemical model and the observations. + Figure 9. allows to investigate the effect of the temperature. density and CO depletion level on the ratio. and to infer under which average conditions our model could reproduce the observed value.," Figure \ref{obs_mod} allows to investigate the effect of the temperature, density and CO depletion level on the ratio, and to infer under which average conditions our model could reproduce the observed value." + It should be noted that the observational ratio also depends on the assumed density. because of the non-LTE effects described in Section 3.3.2..," It should be noted that the observational ratio also depends on the assumed density, because of the non-LTE effects described in Section \ref{outLTE}." +" Agreement between observations and model can only be reached for densities strictly higher than 10? cem"".", Agreement between observations and model can only be reached for densities strictly higher than $^5$ $^{-3}$. + This is consistent with the relatively high measured average density of the core ?, This is consistent with the relatively high measured average density of the core . +) Even at the high density of 105 cem. the model can only reproduce the large observed p-D;:H /o-H;D' ratio if 7210 KK and if the CO depletion level is substantial.," Even at the high density of $^6$ $^{-3}$ , the model can only reproduce the large observed $_2$ $^+$ $_2$ $^+$ ratio if $T>$ K and if the CO depletion level is substantial." + The model points to the fact that the CO abundance should be « 107., The model points to the fact that the CO abundance should be $<$ $^{-5}$. + This corresponds to a CO depletion level of more than 10 (for 722 KK)., This corresponds to a CO depletion level of more than 10 (for $T$ K). + The required depletior is even more severe if the density Is a bit lower. or if the temperature is closer to 10KK. Complete depletion (1.6. >> 100) is required at LOKK. have studied the level of CO depletion toward a sample of seven prestellar cores.," The required depletion is even more severe if the density is a bit lower, or if the temperature is closer to K. Complete depletion (i.e. $>>$ 100) is required at K. have studied the level of CO depletion toward a sample of seven prestellar cores." + These cores are located at distances between 120 and ppc. and three of them are in the Ophiuchus molecular cloud.," These cores are located at distances between 120 and pc, and three of them are in the Ophiuchus molecular cloud." +" The cores have mean densities in the range x 110° cem"". similar to H-MMI."," The cores have mean densities in the range $\times$ $^5$ $^{-3}$, similar to H-MM1." +" Based on IRAM observations of dust continuum and CO isotopologues (at angular resolution of 22""). the authors find depletion factors at the continuum peak of the cores in the range 4.5—-15.5. but this factor drops to 2—5 at offset positions where the continuum is of the peak value."," Based on IRAM observations of dust continuum and CO isotopologues (at angular resolution of $''$ ), the authors find depletion factors at the continuum peak of the cores in the range $-$ 15.5, but this factor drops to $-$ 5 at offset positions where the continuum is of the peak value." + Although the study of Baemann et al. (, Although the study of Bacmann et al. ( +"2002) was doe at a 22"" resolution (to be compared to our 14” beams). preventing us from doing an accurate comparison. the average depletion level we deduce in H-MM1 seems to be in the high range.","2002) was done at a $''$ resolution (to be compared to our $''$ beams), preventing us from doing an accurate comparison, the average depletion level we deduce in H-MM1 seems to be in the high range." +" Moreover the D;:H emission is extended over 40"".", Moreover the $_2$ $^+$ emission is extended over $''$. + Should the H»:D be also extended on that scale. this would imply that the model requires a high CO depletion (>> 10) over an extended region. in contradiction with the results of?.," Should the $_2$ $^+$ be also extended on that scale, this would imply that the model requires a high CO depletion $>>$ 10) over an extended region, in contradiction with the results of." +. It is clear that the observation of the H:D™ spatial distribution and the direct measurement of the level of CO depletion are required to confirm if the astrochemical model ts indeed unable to describe satisfactorily the roots of deuterium chemistry., It is clear that the observation of the $_2$ $^+$ spatial distribution and the direct measurement of the level of CO depletion are required to confirm if the astrochemical model is indeed unable to describe satisfactorily the roots of deuterium chemistry. + Following the discussion on the linewidths in Section 3.1.. the level of the thick observational curves in Fig.," Following the discussion on the linewidths in Section \ref{linewidth}, the level of the thick observational curves in Fig." + 9. is in fact likely underestimated., \ref{obs_mod} is in fact likely underestimated. + Similarly. following the discussion in Section 4.2.. the model predictions are overestimated if the core has not yet reached chemical equilibrium.," Similarly, following the discussion in Section \ref{timescales}, the model predictions are overestimated if the core has not yet reached chemical equilibrium." + These remarks show that the difficulties of the model might be even more severe than shown on Fig. 9.., These remarks show that the difficulties of the model might be even more severe than shown on Fig. \ref{obs_mod}. + The model also seems to exclude kinetic temperatures below KK. but several very cold prestellar cores are known ???).," The model also seems to exclude kinetic temperatures below K, but several very cold prestellar cores are known ." +. The temperature in H-MMIstill needs to bedetermined. but this point could be another difficulty of the model.," The temperature in H-MM1still needs to bedetermined, but this point could be another difficulty of the model." + Further tests of our understanding of the root of deuterium fractionation will depend on the observation of H»D™ , Further tests of our understanding of the root of deuterium fractionation will depend on the observation of $_2$ +"we fitted two-component radial power laws to the observed pre-CME pB radial profiles, from which we obtained a set of coronal electron density profiles all over the region of shock propagation.","we fitted two-component radial power laws to the observed pre-CME $pB$ radial profiles, from which we obtained a set of coronal electron density profiles all over the region of shock propagation." +" In order to determine the location of the shock front, identified as a weak white-light intensity increase located above the expanding CME front (Figure 1, lines), we subtracted the average pre-CME intensity from each frame (see also Ontiveros Vourlidas 2009 for a discussion on white-light shock identification)."," In order to determine the location of the shock front, identified as a weak white–light intensity increase located above the expanding CME front (Figure 1, ), we subtracted the average pre-CME intensity from each frame (see also Ontiveros Vourlidas 2009 for a discussion on white–light shock identification)." +" For each pixel along the front, we extracted the unpolarized white-light brightness of the shocked plasma, Dpost, while the last image acquired before the CME arrival provided the pre-shock brightness, pre, at the same locations in the corona."," For each pixel along the front, we extracted the unpolarized white–light brightness of the shocked plasma, $b_{\rm post}$, while the last image acquired before the CME arrival provided the pre-shock brightness, $b_{\rm pre}$, at the same locations in the corona." +" In each pixel, the brightness b(Z) observed at a projected altitude Z is given by a (LOS) integration of the inferred electron density profile, Πε, multiplied by a geometrical function."," In each pixel, the brightness $b(\zeta)$ observed at a projected altitude $\zeta$ is given by a line-of-sight (LOS) integration of the inferred electron density profile, $n_e$, multiplied by a geometrical function." +" The integration along the LOS is then divided into two integrals: one performed over the unshocked corona (withdensity πο) and the other over a length L across the shocked plasma with density X:n,.", The integration along the LOS is then divided into two integrals: one performed over the unshocked corona (withdensity $n_e$ ) and the other over a length $L$ across the shocked plasma with density $X\cdot n_e$. +" Here, X=pa/Pu is the unknown shock compression ratio and pq (ρι) is the downstream (upstream) plasma density."," Here, $X\equiv\rho_{\rm d}/\rho_{\rm u}$ is the unknown shock compression ratio and $\rho_{\rm d}$ $\rho_{\rm u}$ ) is the downstream (upstream) plasma density." +" In this work L has been estimated as in Bemporad Mancuso (2010), i.e. by assuming that the shock surface has the 3-D shape of an hemispherical shell with thickness equal to that observed on the plane of the sky, corrected for the shock motion during the LASCO/C2 exposure time."," In this work $L$ has been estimated as in Bemporad Mancuso (2010), i.e. by assuming that the shock surface has the 3-D shape of an hemispherical shell with thickness equal to that observed on the plane of the sky, corrected for the shock motion during the LASCO/C2 exposure time." +" Resulting values are L=0.28 Ro and L=0.61 Ro for the shock observed at 11:26 and 11:50 UT, respectively."," Resulting values are $L=0.28$ $_\odot$ and $L=0.61$ $_\odot$ for the shock observed at 11:26 and 11:50 UT, respectively." +" Given L, by adopting the radial density profile derived from the analysis of the pB, the shock compression ratio X can be inferred from the observed Dpre and Dpost."," Given $L$, by adopting the radial density profile derived from the analysis of the $pB$, the shock compression ratio $X$ can be inferred from the observed $b_{\rm pre}$ and $b_{\rm post}$ ." +" The so-obtained compression ratios X along the shock fronts at 11:26 and 11:50 UT, are shown in"," The so-obtained compression ratios $X$ along the shock fronts at 11:26 and 11:50 UT, are shown in" +For a standard accretion disk without a magnetic coupling. depending on the spin of the black hole. among the enerev radiated by the disk up to 6% is captured by the black hole. up to 28% returns to the disk.,"For a standard accretion disk without a magnetic coupling, depending on the spin of the black hole, among the energy radiated by the disk up to $6\%$ is captured by the black hole, up to $28\%$ returns to the disk." + For a disk with a magnetic coupling. more energy is captured by the black hole or returns to the disk.," For a disk with a magnetic coupling, more energy is captured by the black hole or returns to the disk." +" In an extreme case when the magnetic field. lines touch (he disk at the inner boundary. and there is no accretion. depending on the spin of the black hole. among the energy radiated by the disk up to 15% is captured by the black hole. up to 58% returns to the disk. as in the ""infinite efficiency limit” case of a disk magnetically coupled to the material in the transition region (AgolandIxrolik2000)."," In an extreme case when the magnetic field lines touch the disk at the inner boundary and there is no accretion, depending on the spin of the black hole, among the energy radiated by the disk up to $15\%$ is captured by the black hole, up to $58\%$ returns to the disk, as in the “infinite efficiency limit” case of a disk magnetically coupled to the material in the transition region \citep{ago00}." +. Considering the effect of radiation capture. (he radiation flux at the inner part of the disk is moderately mocdified.," Considering the effect of radiation capture, the radiation flux at the inner part of the disk is moderately modified." + IHLowever. in the case of a non-accretion disk magnetically coupled to a Ixerr black hole. the returning radiation dominates al large radii since the radiation due to the magnetic coupling scales as ο while the returning radiation scales as p.* CAgolandIxrolik2000).," However, in the case of a non-accretion disk magnetically coupled to a Kerr black hole, the returning radiation dominates at large radii since the radiation due to the magnetic coupling scales as $r^{-3.5}$ while the returning radiation scales as $r^{-3}$ \citep{ago00}." +. In our calculations the effects of photon capture (either captured by the black hole or returning to the disk) are ienored. we hope to consider them in future.," In our calculations the effects of photon capture (either captured by the black hole or returning to the disk) are ignored, we hope to consider them in future." + For an accretion disk maenetically coupled. to a Ixeir black hole. quasi-steacly state solutions are obtained by assuming the inflow Gme-scale of particles in the disk is much longer than the rotational Uime-scale as adopted in the standard theory of an accretion disk.," For an accretion disk magnetically coupled to a Kerr black hole, quasi-steady state solutions are obtained by assuming the inflow time-scale of particles in the disk is much longer than the rotational time-scale – as adopted in the standard theory of an accretion disk." + Though the magnetic field frozen in the disk slowly moves toward the central black hole as accretion goes on. the inflow velocity of the magnetic field and the particles is much smaller than the rotation velocity of the disk.," Though the magnetic field frozen in the disk slowly moves toward the central black hole as accretion goes on, the inflow velocity of the magnetic field and the particles is much smaller than the rotation velocity of the disk." + Thus. within one rotation period the magnetic field configuration can approximately be regarded as unchanged. ancl the assumption of a state can be applied.," Thus, within one rotation period the magnetic field configuration can approximately be regarded as unchanged, and the assumption of a quasi-steady state can be applied." + For a general distribution of the magnetic field connecting the disk to the black hole. the solutions for the radiation flix aud the internal viscous torque are eiven bv equation (16)) and equation (17)). which are superpositions of the the contribution from accretion and the contribution from magnetic coupling.," For a general distribution of the magnetic field connecting the disk to the black hole, the solutions for the radiation flux and the internal viscous torque are given by equation \ref{flux1}) ) and equation \ref{torque1}) ), which are superpositions of the the contribution from accretion and the contribution from magnetic coupling." + From the view of a long period of time (e.g. with many rotation periods). the radiation flux aud the internal torque of the disk varies with time (the magnetic connection may even disappear finally). but equation (16)) and equation (17)) eive the instant values of the radiation flux ancl the internal torque ol the disk.," From the view of a long period of time (e.g. with many rotation periods), the radiation flux and the internal torque of the disk varies with time (the magnetic connection may even disappear finally), but equation \ref{flux1}) ) and equation \ref{torque1}) ) give the instant values of the radiation flux and the internal torque of the disk." + These general solutions clearly show that lor any distribution of a magnetic field on the disk. the torque produced by the magnetic coupling propagates outward only. (hus (he internal torque and the radiation flux are always zero at the inner boundary of the disk.," These general solutions clearly show that for any distribution of a magnetic field on the disk, the torque produced by the magnetic coupling propagates outward only, thus the internal torque and the radiation flux are always zero at the inner boundary of the disk." +in the line of sight peculiar. velocity. over the scale. of 100 kpe for smaller impact parameters.,in the line of sight peculiar velocity over the scale of 100 kpc for smaller impact parameters. + At such small impact parameters. nearly the full peculiar expansion velocity of the uncderdense region will be probed. while at. larger impact parameters the contribution of the expansion velocity to the line broadening will be reduced by the cecrease in the projection along the line of sight.," At such small impact parameters, nearly the full peculiar expansion velocity of the underdense region will be probed, while at larger impact parameters the contribution of the expansion velocity to the line broadening will be reduced by the decrease in the projection along the line of sight." + After reionization. the Doppler parameters are againo nearly independent of recshilt and impact parameter. suggesting again hat they are," After reionization, the Doppler parameters are again nearly independent of redshift and impact parameter, suggesting again that they are" +show larger decay. timescales. so even higher resolution 3D models (i.e. ο>32) should not be required to finc La lower limit on tlie decay timescale.,"show larger decay timescales, so even higher resolution 3D models (i.e. $N_\phi>32$ ) should not be required to find a lower limit on the decay timescale." + This figureOm also shows that the radiative efficiency of the current laver is crucial to whether the umagnuetosphliere survives for loug periods of time., This figure also shows that the radiative efficiency of the current layer is crucial to whether the magnetosphere survives for long periods of time. + The force-Iree solutious contiuue rapid clissipatiou due to a lack of plasina pressure within the current layer., The force-free solutions continue rapid dissipation due to a lack of plasma pressure within the current layer. + The decay timescale at this resolution for the force-[ree models is 7LOGAL/e? for a=0 and 720GAL/e? for a=0.99.," The decay timescale at this resolution for the force-free models is $\tau\sim +10GM/c^3$ for $a=0$ and $\tau\sim 20GM/c^3$ for $a=0.99$." + The force-Iree simulations show a decay timescale comparable to vacuum «dipole decay on a black hole. which decays as ~(/—19).| starting after only /~20GAL/c?ni (Baumgarte&Shapiro2003).," The force-free simulations show a decay timescale comparable to vacuum dipole decay on a black hole, which decays as $\sim (t-19)^{-4}$ starting after only $t\sim 20GM/c^3$ \citep{bs03}." +. A force-free simulation code with less dissipation may help avoid such fast dissipation iu the lorce-[ree limit (e.g. Spitkovsky 2006)). unless the condition B?E-.," A force-free simulation code with less dissipation may help avoid such fast dissipation in the force-free limit (e.g. \citealt{spit06}) ), unless the condition $B^2E^2$." + This figure also shows that non-zero black hole spin iuduction of poloidal currents cause an extension of ie. timescale for dissipation of the magnetosphere in either the loree-[ree or MHD limits., This figure also shows that non-zero black hole spin induction of poloidal currents cause an extension of the timescale for dissipation of the magnetosphere in either the force-free or MHD limits. + Overall. ie field lines that were initially conuected to the surface of the ((that woul ave reached to iufinity [or a rotating iuagnetosphere) remain counected to the black hole horizon for times much longer than 206À//c* that is predicted by the no-hair theorem.," Overall, the field lines that were initially connected to the surface of the (that would have reached to infinity for a rotating magnetosphere) remain connected to the black hole horizon for times much longer than $\sim 20GM/c^3$ that is predicted by the “no-hair” theorem." + In case of ar initial ccorresponcdiug to the aligned pulsar. those field lines that are associated with the closed part of the uaguetosphere a'e absorbed by thehole.. while those field lines that would reach to iulinity for the uumaenetosphlere are forced to become open by poloidal currents driven into the maguetosphere by the black hole rotation.," In case of an initial corresponding to the aligned pulsar, those field lines that are associated with the closed part of the magnetosphere are absorbed by the, while those field lines that would reach to infinity for the magnetosphere are forced to become open by poloidal currents driven into the magnetosphere by the black hole rotation." + As we showed above through analytical estimates and numerical simulations. the ool a newly [οιmed. rrelaxes to a split-inonopole-tvpe structure.," As we showed above through analytical estimates and numerical simulations, the of a newly formed relaxes to a split-monopole-type structure." + The resulting current sheet is subject to resistive dissipation that would reconnect the field lines from te differeut heiuispheres. producing a set ol closed field lines that will be quickly absorbed by the aaud a set of open field lines that will be released to infiui vin anevent qualitatively similar to solar coronal Wass ejectious (CMEs) (Aschwanden2005)..," The resulting current sheet is subject to resistive dissipation that would reconnect the field lines from the different hemispheres, producing a set of closed field lines that will be quickly absorbed by the and a set of open field lines that will be released to infinity in an event qualitatively similar to solar coronal mass ejections (CMEs) \citep[][]{2005psci.book.....A}." + This will lead to a decrease of the uumber of magnetic flux tubes through each hemisphere Vy., This will lead to a decrease of the number of magnetic flux tubes through each hemisphere $N_B$. + The wwill be slowly balding., The will be slowly balding. +" This ""hair loss"" will proceed o tlie resistive time scale of tlie equatorial current sheet.", This “hair loss” will proceed on the resistive time scale of the equatorial current sheet. + Typically. resistive time scales iui a plasma are much louger than tlie dyuamical times," Typically, resistive time scales in a plasma are much longer than the dynamical times" +"emission at 244m aand 100,m sshow pronounced linear correlations with the eemission within the rregion.","emission at $\,\mu$ and $\,\mu$ show pronounced linear correlations with the emission within the region." +" In the region C, the intensities of the eemission and the continuum emission at 24 and μπι rremain almost constant."," In the region $C$, the intensities of the emission and the continuum emission at 24 and $\,\mu$ remain almost constant." + The CO(2-1) intensity in the rregion (A) is only poorly correlated with the eemission., The CO(2–1) intensity in the region $A$ ) is only poorly correlated with the emission. + In regions B and C the CO(2-1) intensity shows no correlation with the, In regions $B$ and $C$ the CO(2–1) intensity shows no correlation with the +beamed neutron star.,beamed neutron star. + Phe more recent discoveries of strong. broad. ο LH] emission lines (Zepf et al.," The more recent discoveries of strong, broad [O III] emission lines (Zepf et al." + 2007:2008). from this svstem ruled out the possibilities of substantial beaming (since the Ο LL] forbidden line emission. must come from an optically thin and hence isotropically emitting region — see e.g CGnedin et al.," 2007;2008) from this system ruled out the possibilities of substantial beaming (since the [O III] forbidden line emission, must come from an optically thin and hence isotropically emitting region – see e.g Gnedin et al." + 2009)., 2009). +" ""Ehis result thus strongly. supports a scenario with a Lye radio selection produces some biased subsample.," One fundamental issue is whether radio–loud AGN are representative of all AGN, in which case the results derived in this paper would be applicable to all AGN samples, or whether the radio selection produces some biased subsample." + Until it is ully understood what causes an AGN to become radio—loud. this question cannot be definitively answered.," Until it is fully understood what causes an AGN to become radio–loud, this question cannot be definitively answered." + However. studies of ye most powerful radio-loud quasars have found no significant differences between either the host galaxies or the environments of these objects and those of radio-quiet quasars selected to have qe same optical luminosities (e.g. McLure 11999: Nolan 22001: MeLure Dunlop 2001).. except that radio-loud quasars are limited to the absolute upper end of the black hole mass function (Alayc107 ALL. whilst the radio quiet quasars extend to slightly lower values (AJ75107AZ.: eg. Dunlop 22003).," However, studies of the most powerful radio–loud quasars have found no significant differences between either the host galaxies or the environments of these objects and those of radio–quiet quasars selected to have the same optical luminosities (e.g. McLure 1999; Nolan 2001; McLure Dunlop \nocite{mcl99,nol01,mcl01}, except that radio–loud quasars are limited to the absolute upper end of the black hole mass function $M_{\rm +bh} > 10^9 M_{\odot}$ ), whilst the radio quiet quasars extend to slightly lower values $M_{\rm bh} > 5 \times 10^8 M_{\odot}$; e.g. Dunlop 2003)." + This suggests that. at least if analysis is restricted to the most powerful AGN. the radio-loud AGN are likely to be a reasonably unbiased subsample.," This suggests that, at least if analysis is restricted to the most powerful AGN, the radio–loud AGN are likely to be a reasonably unbiased subsample." + At lower optical and radio luminosities the situation is less clear., At lower optical and radio luminosities the situation is less clear. + Given that the AGN fraction seems to depend more strongly on the large-scale than the local environment of their host galaxy. it will be interesting to investigate the dependence upon still larger scale.," Given that the AGN fraction seems to depend more strongly on the large–scale than the local environment of their host galaxy, it will be interesting to investigate the dependence upon still larger scale." + A visual comparison of the AGN and galaxy ccluster distributions suggests that the AGN trace the large-scale structure of the groups and clusters. whilst avoiding the centres of the richest clusters: the very patchy redshift coverage of the 24FGRS at the time that the θα radio sample was constructed makes a quantitative analysis infeasible. however.," A visual comparison of the AGN and galaxy cluster distributions suggests that the AGN trace the large–scale structure of the groups and clusters, whilst avoiding the centres of the richest clusters; the very patchy redshift coverage of the 2dFGRS at the time that the Sad02 radio sample was constructed makes a quantitative analysis infeasible, however." +" Qualitativelv. it is interesting to note there are already 9 radio-loud AGN associated with the largest supercluster structure in the 2dFGRS (at RA 13h. Dec -27. 2~0.084. spanning a physical size of about MMpc east-west), and the very incomplete redshift coverage for this supercluster when the SadO02 analysis took place suggests that this number will increase significantly."," Qualitatively, it is interesting to note there are already 9 radio–loud AGN associated with the largest supercluster structure in the 2dFGRS (at RA 13h, Dec $^{\circ}$, $z \sim 0.084$, spanning a physical size of about Mpc east–west), and the very incomplete redshift coverage for this supercluster when the Sad02 analysis took place suggests that this number will increase significantly." + This is reminiscent of the recent results of Brand shorteitebra03:: these authors used the three-dimensional distribution of NVSS radio sources. as determined through a spectroscopic follow-up. programme in dedicated sky regions. to discover at least one. and possibly two. MMpe-scale superstructures at 20.3 through radio source overdensities.," This is reminiscent of the recent results of Brand \\shortcite{bra03}: these authors used the three–dimensional distribution of NVSS radio sources, as determined through a spectroscopic follow–up programme in dedicated sky regions, to discover at least one, and possibly two, Mpc–scale superstructures at $z \sim 0.3$ through radio source overdensities." + These results suggest that studies of overdensities of radio sources in deep surveys may prove to be a powerful tool for discovering high redshift superclusters., These results suggest that studies of overdensities of radio sources in deep surveys may prove to be a powerful tool for discovering high redshift superclusters. + In line with the argument above that radio-loud AGN are not a highly biased sample of AGN. it is not only radio sources which could be used for this: optical AGN overdensities may also trace these large structures (e.g. Williger 22002: Haines and especially the peripheral regions of these (e.g. Sócchting One of the most striking results to come out of the 2dFGRS analvses is the difference between the properties of the emission and absorption line AGN.," In line with the argument above that radio–loud AGN are not a highly biased sample of AGN, it is not only radio sources which could be used for this: optical AGN overdensities may also trace these large structures (e.g. Williger 2002; Haines \nocite{wil02b,hai03} and especially the peripheral regions of these (e.g. Söcchting \nocite{soc02} + One of the most striking results to come out of the 2dFGRS analyses is the difference between the properties of the emission and absorption line AGN." + Despite having comparable host galaxy properties and covering a similar range of radio luminosities. the environments of these two AGN classes differ dramatically.," Despite having comparable host galaxy properties and covering a similar range of radio luminosities, the environments of these two AGN classes differ dramatically." + In low density environments there is a roughly equal split between absorption and emission line AGN. but in the densest environments essentially none of the AGN show strong line emission.," In low density environments there is a roughly equal split between absorption and emission line AGN, but in the densest environments essentially none of the AGN show strong line emission." +" This is reminiscent of the result of X-ray studies of nearby clusters. which found that the majority of the cluster X-ray AGN were associated with passive elliptical galaxies (?) The difference between emission and absorption-line AGN may be related to that between “normal. strong emission line’ radio galaxies and the class of ""weak emission line’ or ""low excitation radio galaxies found within high radio power samples (2:?).."," This is reminiscent of the result of X-ray studies of nearby clusters, which found that the majority of the cluster X–ray AGN were associated with passive elliptical galaxies \cite{mar02a} + The difference between emission and absorption–line AGN may be related to that between `normal, strong emission line' radio galaxies and the class of `weak emission line' or `low excitation' radio galaxies found within high radio power samples \cite{hin79,lai94}." + In recent years it has been argued that low-excitation radio galaxies do not partake in the orientation-based unified schemes for radio sources (e.g. Barthel 1989).. in which their stronger emission line counterparts are believed to be drawn from the same parent population as radio-loud quasars. but oriented such that direct quasar light is obscured by a torus of dusty material partially surrounding the nucleus.," In recent years it has been argued that low-excitation radio galaxies do not partake in the orientation–based unified schemes for radio sources (e.g. Barthel \nocite{bar89}, in which their stronger emission line counterparts are believed to be drawn from the same parent population as radio–loud quasars, but oriented such that direct quasar light is obscured by a torus of dusty material partially surrounding the nucleus." + Harvanek shorteiteharOl recently used galaxy number count analyses to show that radio galaxies at 2~0.3 typically live in richer environments than radio-loud quasars. and suggested that this might pose a problem for unified schemes.," Harvanek \\shortcite{har01} recently used galaxy number count analyses to show that radio galaxies at $z \sim 0.3$ typically live in richer environments than radio–loud quasars, and suggested that this might pose a problem for unified schemes." + Hardcastle (2) subsequently showed that this was not the case: if the low excitation radio galaxies are removed from the analysis. then there is no significant difference between the environments of strong emission line radio galaxies and those of the radio loud quasars.," Hardcastle \shortcite{har03} subsequently showed that this was not the case: if the low excitation radio galaxies are removed from the analysis, then there is no significant difference between the environments of strong emission line radio galaxies and those of the radio loud quasars." + The entire difference in the Harvanek shorteiteharOl— study. was driven by the low excitation radio galaxies. essentially all of which are found to reside in much richer environments than those of the strong emission line radio galaxies and quasars.," The entire difference in the Harvanek \\shortcite{har01} study was driven by the low excitation radio galaxies, essentially all of which are found to reside in much richer environments than those of the strong emission line radio galaxies and quasars." + This result exactly mirrors that derived in this paper. and suggests that these low excitation radio sources are simply higher radio power examples of the radio sources classified here as absorption-line AGN: the correlation between radio power and environmental richness found in this paper naturally explains why these high radio power low excitation radio sources typically live in very rich environments.," This result exactly mirrors that derived in this paper, and suggests that these low excitation radio sources are simply higher radio power examples of the radio sources classified here as absorption–line AGN; the correlation between radio power and environmental richness found in this paper naturally explains why these high radio power low excitation radio sources typically live in very rich environments." + A key question is what drives this difference between the emission and absorption line AGN?, A key question is what drives this difference between the emission and absorption line AGN? + Are these fundamentally different types of AGN. or does the surrounding environment influence the host galaxy to such an extent that different AGN properties are seen?," Are these fundamentally different types of AGN, or does the surrounding environment influence the host galaxy to such an extent that different AGN properties are seen?" + Although fundamental differences in the AGN properties cannot be ruled out. the similarity of the host galaxies of the two different classes makes this hypothesis difficult to explain. whilst the alternative might be plausible through one or both of two different mechanisms.," Although fundamental differences in the AGN properties cannot be ruled out, the similarity of the host galaxies of the two different classes makes this hypothesis difficult to explain, whilst the alternative might be plausible through one or both of two different mechanisms." + Rawlings Saunders (2) showed that for the strong emission line radio sources. there is a tight correlation between the jet power of the radio source and the emission line luminosity.," Rawlings Saunders \shortcite{raw91b} showed that for the strong emission line radio sources, there is a tight correlation between the jet power of the radio source and the emission line luminosity." + Barthel, Barthel +which the ealactic disc is built. through the SER.,"which the galactic disc is built, through the SFR." + This is well illustrated in mocels reviewed by Losi (1996)., This is well illustrated in models reviewed by Tosi (1996). + Because the infall rate is usually taken as a decreasing function of time. the SER history follows.," Because the infall rate is usually taken as a decreasing function of time, the SFR history follows." + Thus. since infall is so widely used. to explain the lack of metal-poor dwarls observed in the solar neighbourhood. a desirable test of these models would be the age distribution of dwarfs that make the metallicity distribution.," Thus, since infall is so widely used to explain the lack of metal-poor dwarfs observed in the solar neighbourhood, a desirable test of these models would be the age distribution of dwarfs that make the metallicity distribution." + Unfortunately. as mentioned previously. most of the stars in the sample have a main sequence evolution that is too close to the ZAMS to have an age determined with some accuracy in the LER diagram.," Unfortunately, as mentioned previously, most of the stars in the sample have a main sequence evolution that is too close to the ZAMS to have an age determined with some accuracy in the HR diagram." + llowever. an interesting information is available for the voung cbwarfs because they can be detected as X-ray emitters due to coronal activity.," However, an interesting information is available for the young dwarfs because they can be detected as X-ray emitters due to coronal activity." + The cross-icentification between the IOSAT survey and the Catalogue of Nearby Stars (Lluenschetal.1999). can be used to quantify the percentage of potentially voung stars in our sample., The cross-identification between the $ROSAT$ survey and the Catalogue of Nearby Stars \cite{99HUE_EA} can be used to quantify the percentage of potentially young stars in our sample. + There are 90 stars (40 per cent) in our sample in the list of objects resulting from this crosscidentification., There are 90 stars (40 per cent) in our sample in the list of objects resulting from this cross-identification. + Fig. 9((, Fig. \ref{nainesxbv.ps}( ( +a) shows the X-ray [lux ofthese stars as a function of D.V colour.,a) shows the X-ray flux of these stars as a function of $B-V$ colour. + The figure shows that our sample may contain slightly over 90 stars with X-ray. luminosity brighter than log N/bpw=-5.5. which could be absent from the sample of Lluensch et al. (," The figure shows that our sample may contain slightly over 90 stars with X-ray luminosity brighter than log $_{bol}$ =-5.5, which could be absent from the sample of Huensch et al. (" +1999). due to the incompleteness to the X-ray lata for the reddest stars.,"1999), due to the incompleteness to the X-ray data for the reddest stars." + In order to get a rough estimate 10 Corresponding ages. we utilise the log Lx Γιος gx relation of Sterzik Schmitt (1997) and the g loglt«-lost relation of Socderblom ct al. (," In order to get a rough estimate the corresponding ages, we utilise the log $_X$ $_{bol}$ $_{HK}$ relation of Sterzik Schmitt (1997) and the $_{HK}$ -logt relation of Soderblom et al. (" +1991).,1991). +" According to. these relations. log Lxy/Lí,2-5.5 corresponds to an age of 2 Gyr."," According to these relations, log $_X$ $_{bol}$ =-5.5 corresponds to an age of 2 Gyr." + Most our stars having X-ray emission at the level of log Lx/Li 2-5.5A could therefore be considered: vounger than 2 Gvr., Most our stars having X-ray emission at the level of log $_X$ $_{bol}>$ -5.5 could therefore be considered younger than 2 Gyr. + Fig. 9((, Fig. \ref{nainesxbv.ps}( ( +b) illustrates what the metallicity clistribution of these stars ds.,b) illustrates what the metallicity distribution of these stars is. + 1t is centred. on Fe/1I]20.0.. with an unexpected but. important contribution of 11 stars with Fefl]e-0.3.," It is centred on [Fe/H]=0.0, with an unexpected but important contribution of 11 stars with $<$ -0.3." + Some of these stars have a spectroscopic iron abundance which confirms the photometric iron abundance., Some of these stars have a spectroscopic iron abundance which confirms the photometric iron abundance. + For instance. HP ss622 has ως 0.506. and ρωσ. dated 15 Gyr in (Exlvardssonetal. 1902). and. LLP 15510 has pir. -0.35. and epecro=-O4S8 (Pasquinietal.1994).," For instance, HIP 88622 has $_{photo}$ =-0.506, and $_{spectro}$ =-0.47, dated 15 Gyr in \shortcite{93EDV_EA2}, and HIP 15510 has $_{photo}$ =-0.35, and $_{spectro}$ =-0.48 \cite{94PAS_EA}." + The X-ray emission. of LIP 88622. may appear somewhat puzzling recause of its metallicity. and. because HIP. 88622 has no significant. chromospheric emission. according to Pasquini et al. (," The X-ray emission of HIP 88622 may appear somewhat puzzling because of its metallicity, and because HIP 88622 has no significant chromospheric emission according to Pasquini et al. (" +1994).,1994). + While these cases may be exceptions. it has oen suggested. more eenerally that activity may allect whotometric indices and that presumably voung stars may ος measured. as delicient objects if abundance is measured rom photometry (Gimenezetal.1991:Moraleal...1996:Favataetal..1997:Rocha-Pinto&Maciel.," While these cases may be exceptions, it has been suggested more generally that activity may affect photometric indices and that presumably young stars may be measured as deficient objects if abundance is measured from photometry \cite{91GIM_EA,96MOR_EA,97FAV_EA3,98ROC_EA}." + 1998).. This effect may explain the shape of the histogram of Fig.9.. but its importance is uncertain. and we do not try to correct it.," This effect may explain the shape of the histogram of \ref{nainesxbv.ps}, but its importance is uncertain, and we do not try to correct it." + While some of the stars may have been detected: as X-ray emitters in the sample for reasons not related to their age. it is clearly demonstrated. that X-ray. enission is à tracer of voung stars (Cruilloutctabl.1998).," While some of the stars may have been detected as X-ray emitters in the sample for reasons not related to their age, it is clearly demonstrated that X-ray emission is a tracer of young stars \cite{98GUI_EA}." +. In this respect. the [act that around 40 per cent of the sample is composed of X-ray emitters that have ages less than 2 Civr suggests a [ew comments.," In this respect, the fact that around 40 per cent of the sample is composed of X-ray emitters that have ages less than 2 Gyr suggests a few comments." + In section 77.. we calculate. a model metallicity istribution assuming a constant SER.," In section \ref{models}, we calculate a model metallicity distribution assuming a constant SFR." + We argue that dthough a constant SER is not phenomenologically directly related to the gas content. it is compatible with available determination of the SER history in the Milky Was.," We argue that although a constant SFR is not phenomenologically directly related to the gas content, it is compatible with available determination of the SFR history in the Milky Way." + Loa constant star formation rate has dominated. the history of 1e galactic disc. then of the order of 1525 per cent of the stars are expected to have an age less than 2 Civr (for a 8-12 Car thin disc).," If a constant star formation rate has dominated the history of the galactic disc, then of the order of 15–25 per cent of the stars are expected to have an age less than 2 Gyr (for a 8-12 Gyr thin disc)." + However. a distance limited. sample of the solar. neighbourhood is biased against old stars because of," However, a distance limited sample of the solar neighbourhood is biased against old stars because of" +"The suite of simulations described above shows that over an extremely wide range of parameters, the ultimate consequence of the interaction between a galactic outflow and a primordial minihalo is the rapid formation of dense clusters containing up to a few 10°Mo of stars.","The suite of simulations described above shows that over an extremely wide range of parameters, the ultimate consequence of the interaction between a galactic outflow and a primordial minihalo is the rapid formation of dense clusters containing up to a few $10^6 M_\odot$ of stars." +" While such high-redshift clusters are not directly observable with current telescopes, their rapid bursts of star formation and consequent low mass-to-light ratios present a opportunity for study with the next generation of instruments."," While such high-redshift clusters are not directly observable with current telescopes, their rapid bursts of star formation and consequent low mass-to-light ratios present a opportunity for study with the next generation of instruments." +" At the same time, their compact nature and formation in low-density environments makes it likely that many of them may have survived to the present, allowing for indirect connections with current stellar populations."," At the same time, their compact nature and formation in low-density environments makes it likely that many of them may have survived to the present, allowing for indirect connections with current stellar populations." + Here we explore both of these connections., Here we explore both of these connections. +" To calculate the direct observability of the clusters formed in our simulations, we first constructed an estimate of the number of stars formed as a function and time and position."," To calculate the direct observability of the clusters formed in our simulations, we first constructed an estimate of the number of stars formed as a function and time and position." +" As there is no explicit prescription for star formation included in our simulations, we instead built up the star-formation history by post-processing our outputs, carrying out the following steps:"," As there is no explicit prescription for star formation included in our simulations, we instead built up the star-formation history by post-processing our outputs, carrying out the following steps:" +Array (CTA) and the conditions that must be fulfilled to achieve detection.,Array (CTA) and the conditions that must be fulfilled to achieve detection. + We carried out continuum observations with the Very Large Array (NRAO!)) in two array configurations: C at 1.42 GHz in April 2008 and D at 4.86 GHz in August 2008., We carried out continuum observations with the Very Large Array ) in two array configurations: C at 1.42 GHz in April 2008 and D at 4.86 GHz in August 2008. +" The set up allowed mapping largest structures of the size of the MSX source and also guaranteed good matching beams of 12"".", The set up allowed mapping largest structures of the size of the MSX source and also guaranteed good matching beams of $12''$. + The flux calibrator used was01374-331?., The flux calibrator used was. +. Phase calibrator scans —2052+365 at L band and 20074-404 at C band- were interleaved with target scans., Phase calibrator scans –2052+365 at L band and 2007+404 at C band– were interleaved with target scans. + The total time on source was 3h at each band and 50 MHz for the total bandwidth., The total time on source was 3h at each band and 50 MHz for the total bandwidth. + The data was calibrated with the AIPS package in the standard way and analyzed with the Miriad routines., The data was calibrated with the AIPS package in the standard way and analyzed with the Miriad routines. + We used to produce robust-weighted images., We used to produce robust-weighted images. +" Figure 2 presents the resulting images after primary beam correction, re-gridded with the same synthesized beam."," Figure 2 presents the resulting images after primary beam correction, re-gridded with the same synthesized beam." + At both frequencies there is emission along the extension of the MSX source., At both frequencies there is emission along the extension of the MSX source. +" The radio source is larger at 1.42 GHz, toward the eastern half of IR contours of the bow shock (increasing right ascensions)."," The radio source is larger at 1.42 GHz, toward the eastern half of IR contours of the bow shock (increasing right ascensions)." +" Clearly, the spectral index a (S,ος v?) is negative in there."," Clearly, the spectral index $\alpha$ $S_{\nu} \propto \nu^{\alpha}$ ) is negative in there." +" In both images, the rms attained is similar (0.3 mJy beam-! at 1.42 GHz and 0.2 at 4.8, 12"" synthesized beam)."," In both images, the rms attained is similar (0.3 mJy $^{-1}$ at 1.42 GHz and 0.2 at 4.8, $12''$ synthesized beam)." +" A detached ellipsoidal source at a, 6[J2000] = 20h 33™35®,44?04/30"", called ES here, is thestrongest one in the field."," A detached ellipsoidal source at $\alpha,\delta$ [J2000] = $20^{\rm h}$ $33^{\rm m}35^{\rm s}, 44^{\circ} 04' 30''$, called ES here, is thestrongest one in the field." + The fluxes are Spg(1.4GHz)=105+10 mJy and Sgs(4.8GHz)=95+5 mJy., The fluxes are $S_{\rm ES} (\rm 1.4GHz)= 105\pm10$ mJy and $S_{\rm ES}(\rm 4.8GHz)= 95\pm 5$ mJy. + ES spectral index =—0.1€0.1 is characteristic of an optically thin region.," ES spectral index $<\alpha> = +-0.1\pm0.1$ is characteristic of an optically thin region." +" 'The hypothesis of a physical association between the star and the radio/IR features is supported by the very good agreement of the residual proper motion of the star (15.7°+9.4° east of north, see Comerónn Pasquali 2007) and the direction from the star to the apsis of the bow shock (8.8?+10? east of north)."," The hypothesis of a physical association between the star and the radio/IR features is supported by the very good agreement of the residual proper motion of the star $15.7^\circ\pm 9.4^\circ$ east of north, see Comerónn Pasquali 2007) and the direction from the star to the apsis of the bow shock $8.8^\circ\pm 10^\circ$ east of north)." + The star velocity vector on the plane of the sky is represented in the upper panel of Fig., The star velocity vector on the plane of the sky is represented in the upper panel of Fig. + 2., 2. + We used the continuum images at 1.42 and 4.86 GHz to build a spectral index distribution map., We used the continuum images at 1.42 and 4.86 GHz to build a spectral index distribution map. + We only considered input pixels with a signal-to-noise ratio > 4., We only considered input pixels with a signal-to-noise ratio $\geq$ 4. +" Besides this, the spectral index map was masked for a signal-to-noise ratio > 10."," Besides this, the spectral index map was masked for a signal-to-noise ratio $\geq$ 10." + Figure 3 shows the spectral index distribution and corresponding noise maps., Figure 3 shows the spectral index distribution and corresponding noise maps. + In Figure 3 we have represented the spectral index values derived from 1.42 and 4.86 GHz data at positions where the radio continuum signals are well above the noise (4 times or more) and the values of the spectral index, In Figure 3 we have represented the spectral index values derived from 1.42 and 4.86 GHz data at positions where the radio continuum signals are well above the noise (4 times or more) and the values of the spectral index +"(also reported in Tab. 1)),","(also reported in Tab. \ref{infogen}) )," + and the last the highest metallicity value found., and the last the highest metallicity value found. +" With the current observations, we can exclude the presence of stars younger than ~500 Myr in these galaxies (from the absence of an upper main sequence or supergiant stars), but above the TRGB of each galaxy we can see a number of stars that are probably luminous AGB stars."," With the current observations, we can exclude the presence of stars younger than $\sim500$ Myr in these galaxies (from the absence of an upper main sequence or supergiant stars), but above the TRGB of each galaxy we can see a number of stars that are probably luminous AGB stars." +" They are the bright tip of the iceberg of an IAP (1 to 9 Gyr), which we expect to find at this position for metal-poor stellar populations."," They are the bright tip of the iceberg of an IAP $\sim1$ to 9 Gyr), which we expect to find at this position for metal-poor stellar populations." +" For populations more metal-rich than [Fe/H]~—1.0, some old and metal-rich stars may also be found above the TRGB, but our target objects are predominantly metal-poor (see Tab. 1)),"," For populations more metal-rich than $\sim-1.0$, some old and metal-rich stars may also be found above the TRGB, but our target objects are predominantly metal-poor (see Tab. \ref{infogen}) )," + so we conclude that the presence of such stars is not significant in our sample., so we conclude that the presence of such stars is not significant in our sample. +" Unfortunately, the low Galactic latitude of the Centaurus A group (b~ 20°) means that a substantial amount of Galactic foreground stars contaminate the CMDs of our target galaxies."," Unfortunately, the low Galactic latitude of the Centaurus A group $b\sim20^{\circ}$ ) means that a substantial amount of Galactic foreground stars contaminate the CMDs of our target galaxies." + We simulate the expected foreground using both the TRILEGAL models (?) and the Besangoon models (?).., We simulate the expected foreground using both the TRILEGAL models \citep{girardi05} and the Besançoon models \citep{robin03}. + An example of the expected contamination in our CMDs is shown in Fig. 3.., An example of the expected contamination in our CMDs is shown in Fig. \ref{cont}. . + We plot both optical and NIR CMDs for CenA- and overplot the positions of foreground stars predicted by the two models (blue crosses for TRILEGAL and red asterisks for Besangoon).," We plot both optical and NIR CMDs for CenA-dE1, and overplot the positions of foreground stars predicted by the two models (blue crosses for TRILEGAL and red asterisks for Besançoon)." +" We convolve the simulated data with photometric errors and take into account incompleteness effects, and we stress that the simulated Galactic stars are only one random realization of the adopted models."," We convolve the simulated data with photometric errors and take into account incompleteness effects, and we stress that the simulated Galactic stars are only one random realization of the adopted models." +" Just for comparison purposes, we only plot stars simulated by Besangoon models that have masses > 0.15Mo, since the TRILEGAL models do not include them in their computation."," Just for comparison purposes, we only plot stars simulated by Besançoon models that have masses $>0.15$ $_{\odot}$, since the TRILEGAL models do not include them in their computation." +" These excluded few, very low-mass stars would have the effect of extending the Besancoon plotted sequence to slightly redder colors."," These excluded few, very low-mass stars would have the effect of extending the Besançoon plotted sequence to slightly redder colors." +" We can notice that at magnitudes Kgx20.5 and colors Jo—Ko€1 (where only foreground stars are found), the number counts resulting from the models are comparable to each other, and similar to the number of observed stars."," We can notice that at magnitudes $K_0\lesssim20.5$ and colors $J_0-K_0\lesssim1$ (where only foreground stars are found), the number counts resulting from the models are comparable to each other, and similar to the number of observed stars." +" We stress, however, that the models give slightly different results, and in particular the Besangoon model reaches slightly redder colors, but a comparison between the models is beyond the goals of this study."," We stress, however, that the models give slightly different results, and in particular the Besançoon model reaches slightly redder colors, but a comparison between the models is beyond the goals of this study." +" Overall, it is clear from the left panel that the luminous AGB region is the most affected in the optical, and we have no way of determining which stars belong to the dwarf galaxy and which are part of the foreground."," Overall, it is clear from the left panel that the luminous AGB region is the most affected in the optical, and we have no way of determining which stars belong to the dwarf galaxy and which are part of the foreground." + We will come back to the NIR CMD in the following Section., We will come back to the NIR CMD in the following Section. +" In ? we gave a rough estimate of the possible fraction of the IAP for the target early-type dwarfs by using the fuel consumption theorem (originally introduced by ?,, see also ?))."," In \citet{crnojevic10} we gave a rough estimate of the possible fraction of the IAP for the target early-type dwarfs by using the fuel consumption theorem (originally introduced by \citealt{renzini86}, see also \citealt{armand93}) )." +" We point out that the IAP fractions we report throughout this work are intended as numbers of stars (i.e., number of luminous AGBstars to old RGB stars)."," We point out that the IAP fractions we report throughout this work are intended as numbers of stars (i.e., number of luminous AGBstars to old RGB stars)." +" The results were the following: for CenA-dE1 and SGC1319.1-4216 the IAP fraction is ~10%, while for ESO269-066 it is ~15% of the entire population."," The results were the following: for CenA-dE1 and SGC1319.1-4216 the IAP fraction is $\sim10\%$, while for ESO269-066 it is $\sim15\%$ of the entire population." + The NIR CMDs are presented in Fig. 4.., The NIR CMDs are presented in Fig. \ref{cmdnir}. . +" We have corrected the magnitudes for foreground reddening, referring to theNED"," We have corrected the magnitudes for foreground reddening, referring to theNED" +and. orbital parameters are already known. presenting ideal laboratories for investigating accretion physics in active galactic nuclei (AGN).,"and orbital parameters are already known, presenting ideal laboratories for investigating accretion physics in active galactic nuclei (AGN)." + Furthermore. if major mergers of ealaxies trigger luminous AGN activity (c.@.. Sandersctal.αιανα&IIaehnelt2000:Llopkinsetal.2007. 2008)). then the characteristic LAL emission. promptly following the SMBII coalescence may. herald the birth of a quasar (Tanaka.Haiman&Alenou2010).," Furthermore, if major mergers of galaxies trigger luminous AGN activity (e.g., \citealt{Sanders+88, Hernquist89, Carlberg90, BarnesHernquist91, + HernquistMihos95, MihosHernquist96, KauffmannHaehnelt00, + Hopkins+07a, Hopkins+08}) ), then the characteristic EM emission promptly following the SMBH coalescence may herald the birth of a quasar \citep{THM10}." +. To date. theoretical studies of LAL signatures of GW-emitting SAIBIL binaries have largely centred on systems. predicted. to be detectable by future space-based laser interferometers such as (LISA). with total mass 107‘(1|2)1M. and redshifts as high as 210.," To date, theoretical studies of EM signatures of GW-emitting SMBH binaries have largely centred on systems predicted to be detectable by future space-based laser interferometers such as ), with total mass $\sim 10^{5-7}(1+z)^{-1}\Msol$ and redshifts as high as $z\gta 10$." + For the purposes of multi-messenger astronomy. a paramount feature of interferometer-detectors is the precision with they could determine the angular sky position of SALBLL sources (Ixocsisetal.2O0G):: to X1deg (Vecchio2004:Lang& 2008).. or perhaps even to <1 when spin-induced precession (Lang&Hughes2006) or higher-order harmonics (AleWilliamsetal.2010) are included in the analysis of the waveform.," For the purposes of multi-messenger astronomy, a paramount feature of interferometer-detectors is the precision with they could determine the angular sky position of SMBH sources \citep{Kocsis+06}: to $\ltsim 1 \deg$ \citep{Vecchio04, + LH08}, or perhaps even to $\ltsim 1^{\prime}$ when spin-induced precession \citep{LH06} or higher-order harmonics \citep{McWill+10} are included in the analysis of the waveform." + In this paper. we evaluate the prospects of performing multi-messenger using pulsar timing arrays (P'PAs). which aim to detect low-frequencey. (nllz) CWs through precision iming of Galactic millisecond pulsars.," In this paper, we evaluate the prospects of performing multi-messenger using pulsar timing arrays (PTAs), which aim to detect low-frequency (nHz) GWs through precision timing of Galactic millisecond pulsars." + A major goal for *PAs is to detect the stochastic GW background. due to he population of compact SALBLL binaries in our cosmic neighborhood (2X; 2)., A major goal for PTAs is to detect the stochastic GW background due to the population of compact SMBH binaries in our cosmic neighborhood $z\ltsim 2$ ). + Jenetetal.(2005) showed that his requires the timing of some Z20 pulsars to ~100ns orecision over 5 vears. a goal that could. be achieved »w currentlv operational arrays (Manchester:2008:Ver-yestetal.2009.anclrefs.therei," \cite{Jenet+05} showed that this requires the timing of some $\ga 20$ pulsars to $\sim 100 ~{\rm ns}$ precision over 5 years, a goal that could be achieved by currently operational arrays \citep[][and refs. therein]{Manchester08, Verbiest+09}." +n) Recent theoretical »opulation-svnthesis studies (Sesana.Vecchio&Volonteri2009:Sesana&VecchioPOLO:IXocsis2011) have xedieted that PTA observations may also be able to the most massive and/or most nearby SMDII yinawics that stick out above the stochastic The primary factor that. determines the number of detectable individual sources is expected to be the array sensitivity. with theoretical uncertainties such as the binary mass function and orbital evolution alfecting npredictions of the observable population by [actors of unity (Sesanactal. 2011).," Recent theoretical population-synthesis studies \citep{SVV09, SesVec10, KocSes11} + have predicted that PTA observations may also be able to the most massive and/or most nearby SMBH binaries that stick out above the stochastic The primary factor that determines the number of detectable individual sources is expected to be the array sensitivity, with theoretical uncertainties such as the binary mass function and orbital evolution affecting predictions of the observable population by factors of unity \citep{Sesana+11}." +. The systems individuallv detected by P'EXs will cüller [rom targets in several wavs., The systems individually detected by PTAs will differ from targets in several ways. + First. according to the population svnthesis studies. the binaries will have chirp Hlasses B/SAL typically around ο1077AZ...," First, according to the population synthesis studies, the binaries will have chirp masses M, typically around $\Mch\sim10^{8.5}\Msol$." + Above. Alo$, against galactocentric radius, $r$, for M51, including all the regions, but distinguishing by point styles among those found in different large scale features of the galaxy: arms, interarm zones, and the central kpc." +" The outstanding feature of this plot is the single straight line fit (which is an exponential, as the electron density scale is logarithmic), which has the form: where (ne),, with value 1041 cm3 is acentral value of electron density and h is a ?scale length, which takes the value 10+1.0 kpc."," The outstanding feature of this plot is the single straight line fit (which is an exponential, as the electron density scale is logarithmic), which has the form: where $\left_o$ , with value $\pm$ 1 $^{-3}$ is acentral value of electron density and $h$ is a scale length, which takes the value $\pm$ 1.0 kpc." +" We have limited the fit to galactocentric radii within r = 10.4 kpc, because beyond this radius there is an abrupt local rise and fall in (ιο), attributable to the effect of the interacting neighbour galaxy NGC 5195."," We have limited the fit to galactocentric radii within $r$ = 10.4 kpc, because beyond this radius there is an abrupt local rise and fall in $\left$, attributable to the effect of the interacting neighbour galaxy NGC 5195." +" We can make separate estimates for the constants (ne), and h in the arms and in the interarm zones.", We can make separate estimates for the constants $\left_o$ and $h$ in the arms and in the interarm zones. +" For the arms this gives us = 10 cm? and h = 11 kpc, while for the interarm (ne),zones we find (ne), = 8 επι” and h = 11 kpc."," For the arms this gives us $\left_o$ = 10 $^{-3}$ and $h$ = 11 kpc, while for the interarm zones we find $\left_o$ = 8 $^{-3}$, and $h$ = 11 kpc." +" Since Ὁ,we know that in the upper panel of Fig.", Since we know that in the upper panel of Fig. +" 3 much of the apparent scatter is due to the (inverse) dependence of (πιο) on the radii of the individual regions, we next opted to apply a normalized correction factor for this effect to all the data, and replot the figure, yielding the central panel of Fig. 3,,"," \ref{F:densidad_vs_distancia_m51} much of the apparent scatter is due to the (inverse) dependence of $\left$ on the radii of the individual regions, we next opted to apply a normalized correction factor for this effect to all the data, and replot the figure, yielding the central panel of Fig. \ref{F:densidad_vs_distancia_m51}," + in which the (exponential) linear fit is clearly much improved., in which the (exponential) linear fit is clearly much improved. +" To make the fit we have chosen to leave out the zones from 1.4 to 4.6 kpc from the center, and from 10.4 to 11.5 kpc from the center, and now find values of and h of 12 cm? and 9 kpc respectively."," To make the fit we have chosen to leave out the zones from 1.4 to 4.6 kpc from the center, and from 10.4 to 11.5 kpc from the center, and now find values of $\left_o$ and $h$ of 12 $^{-3}$ and 9 kpc respectively." +" It is notable (n),that the scale length found for the neutral atomic hydrogen component of M51 by Tilanus&Allen(1991) is 9.1 kpc, on which we will comment further below."," It is notable that the scale length found for the neutral atomic hydrogen component of M51 by \citet{tilanus91} is 9.1 kpc, on which we will comment further below." + We should point out that the zone between the center of the galaxy and a radius of 1.4 kpc does fit the exponential for the disc as a whole., We should point out that the zone between the center of the galaxy and a radius of 1.4 kpc does fit the exponential for the disc as a whole. +" To take a look at the global behavior we have taken averages of all the regions within annuli of width 200 pc and have plotted, in the lower panel of Fig."," To take a look at the global behavior we have taken averages of all the regions within annuli of width 200 pc and have plotted, in the lower panel of Fig." + 3 these values against galactocentric radius.," \ref{F:densidad_vs_distancia_m51} + these values against galactocentric radius." +" This shows the overall trend in (πο), and we can pick out the general decline with radius, but note a plateau between 1.4 kpc and 4.6 kpc."," This shows the overall trend in $\left$, and we can pick out the general decline with radius, but note a plateau between 1.4 kpc and 4.6 kpc." +" If we omit the points in this interval from those entering the exponential fit, we obtain values for (ne), and h, of 11 cm? and 9 kpc respectively."," If we omit the points in this interval from those entering the exponential fit, we obtain values for $\left_o$ and $h$, of 11 $^{-3}$ and 9 kpc respectively." +" 'To compute the effect of the temperature gradient on the determination of and of its gradient, we used data from Bresolinetal.(πο)(2004) giving the O abundance gradient, and the widely used CLOUDY model (Ferlandetal.1998) to derive the effective temperature gradient, which we then incorporated into the expression and recalculated the electron densities."," To compute the effect of the temperature gradient on the determination of $\left$ and of its gradient, we used data from \citet{bresolin04} + giving the O abundance gradient, and the widely used CLOUDY model \citep{ferland98} + to derive the effective temperature gradient, which we then incorporated into the expression and recalculated the electron densities." +"This introduced a change in the exponent of the relation between (τις) and radius from -0.56 to -0.55, and in the scale length from 9.2 to 9.9 kpc.","This introduced a change in the exponent of the relation between $\left$ and radius from -0.56 to -0.55, and in the scale length from 9.2 to 9.9 kpc." + These are the values represented in Fig. 3.., These are the values represented in Fig. \ref{F:densidad_vs_distancia_m51}.. + Here we should point out the similarity of the lower panel of Fig., Here we should point out the similarity of the lower panel of Fig. + 3 with the radial plot of the azimuthally averaged column density in Tilanus&Allen (1991).," \ref{F:densidad_vs_distancia_m51} + with the radial plot of the azimuthally averaged column density in \citet{tilanus91}." +. General expressions for (πο) are then:, General expressions for $\left$ are then: +Reeion.,Region. + Therefore the study of these sources offers us the unique possibility to probe the jet iu its initial portion., Therefore the study of these sources offers us the unique possibility to probe the jet in its initial portion. + The typical Spectral Energy. Distribution of blazars oreseuts two broad laps. the first peasing in the IR through N-rax baud. the second im the 5-axy doniadn. extending in some sources to the TeV band (Fig.1).," The typical Spectral Energy Distribution of blazars presents two broad humps, the first peaking in the IR through X-ray band, the second in the $\gamma $ -ray domain, extending in some sources to the TeV band (Fig.1)." + The hig1 degree of polarization clearly indicate that the low energy conr2010111 is produced through synchrotron emission., The high degree of polarization clearly indicate that the low energy component is produced through synchrotron emission. +" The second peak is likely prodiced hrough IC scattering of soft photons bv the same electron population resposible or the svucwotron enmission (although other mechanisms have been proposed. see c.g, the review in Sikora Madejski WwX01)."," The second peak is likely produced through IC scattering of soft photons by the same electron population responsible for the synchrotron emission (although other mechanisms have been proposed, see e.g. the review in Sikora Madejski 2001)." + In the model it is assuned hat the sof photon cucrey density is dominated by the svuchrotron ploOuS hemsclves. while the so-called models assune. that soft ghotons coming frou the external cuviromment (disk. DLR) dominate «NOY other possible coutributious.," In the model it is assumed that the soft photon energy density is dominated by the synchrotron photons themselves, while the so-called models assume that soft photons coming from the external environment (disk, BLR) dominate over other possible contributions." + Althouga blazars form a rather inhomogeneous class from the poiit of view of the optical classification. the SEDs prescut a notable trend with the hnuninositv. the so-calledsequence (Fossati et al.," Although blazars form a rather inhomogeneous class from the point of view of the optical classification, the SEDs present a notable trend with the luminosity, the so-called (Fossati et al." + 1995: Donao et al., 1998; Donato et al. + 2001)., 2001). + The sources with high luninosity output (L~107 org ean? 1 sources nadnlv classified as Flat Spectrmu Radio Quasars) have both peaks located at low cuereies. the first in the IR region aud the high euergv peak around 1 MeV. Ou the contrary. at low power the sources (BL Lac objects.," The sources with high luminosity output $L\sim +10^{48}$ erg $^{-2}$ $^{-1}$, sources mainly classified as Flat Spectrum Radio Quasars) have both peaks located at low energies, the first in the IR region and the high energy peak around 1 MeV. On the contrary, at low power the sources (BL Lac objects," +There is no CP violation in the leptonic sector of the Standard Model (SAL) of fandamental particles and interactions.,There is no CP violation in the leptonic sector of the Standard Model (SM) of fundamental particles and interactions. + However. in most extensions of SM. (there can be several CP phases.," However, in most extensions of SM, there can be several CP phases." + In (he simplest (μου generation scenario. there is a Dirac (vpe CP violating phase in (he leptonic mixing matrix.," In the simplest three generation scenario, there is a Dirac type CP violating phase in the leptonic mixing matrix." + However. for Majorana neutrinos. there could be two additional phases (Majorana phases).," However, for Majorana neutrinos, there could be two additional phases (Majorana phases)." + It is possible to work in a parameterizations in which all the three CP violating phases are situated in the chareecd current leptonic mixing matrix., It is possible to work in a parameterizations in which all the three CP violating phases are situated in the charged current leptonic mixing matrix. + Without anv loss of generality. one can work in the flavor basis in which the charged lepton mass matrix is diagonal so that the neutrino mass matrix carries all the information about. CP violation.," Without any loss of generality, one can work in the flavor basis in which the charged lepton mass matrix is diagonal so that the neutrino mass matrix carries all the information about CP violation." + In the flavor basis. (he mass matrix for \lajorana neutrinos contains nine physical parameters including the three mass eigenvalues. thiree mixing angles and the (hree CP-violating phases.," In the flavor basis, the mass matrix for Majorana neutrinos contains nine physical parameters including the three mass eigenvalues, three mixing angles and the three CP-violating phases." +" The (wo squared-mnass. differences (Amqz, and Am.) and the two mixing angles (£45 and 054) have been measured in solar. alimospheric aud reactor experiments."," The two squared-mass differences $\Delta m^2_{12}$ and $\Delta m^2_{13}$ ) and the two mixing angles $\theta_{12}$ and $\theta_{23}$ ) have been measured in solar, atmospheric and reactor experiments." + The third mixing angle 644 ancl the Dirac-tvpe phase ὁ are expected to be measured in (he forthcoming neutrino oscillation experimentis., The third mixing angle $\theta_{13}$ and the Dirac-type CP-violating phase $\delta$ are expected to be measured in the forthcoming neutrino oscillation experiments. + The possible measurement of the effective Majorana mass in neutrinoless double ο decav searches will provide an additional constraint on (he remaining three neutrino parameters viz., The possible measurement of the effective Majorana mass in neutrinoless double $\beta$ decay searches will provide an additional constraint on the remaining three neutrino parameters viz. + (he neutrino mass scale and two Majorana-tvpe CP-violating phases., the neutrino mass scale and two Majorana-type CP-violating phases. + While the neutrino mass scale will be independently. determined by the direct beta decay searches and cosmologicalex observations. the two Majorana phases will not," While the neutrino mass scale will be independently determined by the direct beta decay searches and cosmological observations, the two Majorana phases will not" +Burbidgeetal.(1957). attributed the production of the isotopes heavier (han the iron eroup to three processes of nucleosyntliesis. the r- ancl s-processes of neutron addition. and the p-process of proton addition.,"\citet{Bur57} attributed the production of the isotopes heavier than the iron group to three processes of nucleosynthesis, the $r$ - and $s$ -processes of neutron addition, and the $p$ -process of proton addition." +" The conditions they specified for the p-process. proton densities pX,~LO? g ?oE and temperatures T—23x10? IX. were difficult to realize in nature and so other processes ancl sites were sought."," The conditions they specified for the $p$ -process, proton densities $\rho X_p \sim 10^2$ g $^{-3}$ and temperatures $T \sim +2 - 3 \times 10^9$ K, were difficult to realize in nature and so other processes and sites were sought." + Arnould(1976) and Wooslev(1978). attributed the production of the p-process nuclei to photocdisintegration. a series of (5.n). (5.p) aud (5.0) reactions flowing downward through radioactive proton-rich progenitors Irom lead to iron.," \citet{Arn76} and \citet{Woo78} + attributed the production of the $p$ -process nuclei to photodisintegration, a series of $\gamma,$ n), $\gamma,$ p) and $\gamma,\alpha$ ) reactions flowing downward through radioactive proton-rich progenitors from lead to iron." +" Their 5-process"" operated upon previously existing s-process seed in the star to make the p-process. and was (hus ""secondary in nature (or even ""tertiary"" since (he s-process itself is secondary)."," Their $\gamma$ -process” operated upon previously existing $s$ -process seed in the star to make the $p$ -process, and was thus “secondary” in nature (or even “tertiary” since the $s$ -process itself is secondary)." + It could only happen in a star niade from (he ashes of previous stars that had made the s-process., It could only happen in a star made from the ashes of previous stars that had made the $s$ -process. + Arnould suggested hvdrostatie oxveen burning in massive stars as (he site where the necessary conditions were realized: Woosley and Howard. who discovered the relevant nuclear flows independently. discussed explosive oxveen and neon burning in à Type II supernova as (he likely site.," Arnould suggested hydrostatic oxygen burning in massive stars as the site where the necessary conditions were realized; Woosley and Howard, who discovered the relevant nuclear flows independently, discussed explosive oxygen and neon burning in a Type II supernova as the likely site." + Over the vears. increasingly refined cealeulations showed (hat a portion of the p-nuclei could actually be produced as Wooslev and Llowarcl described," Over the years, increasingly refined calculations showed that a portion of the $p$ -nuclei could actually be produced as Woosley and Howard described" +components is a Πο WD.,components is a He WD. + Double Πο WD οσους will nudereo a very similar evolution as we have described for C/O WD increers. with appropriately adjusted energy and timescales: a We-burning shell will ignite oft-ceuter and diffuse inwards. possibly with episoclic flashes due to the ineficicney of neutrino cooling (ben1990:Saio&Nomoto 1998).," Double He WD mergers will undergo a very similar evolution as we have described for C/O WD mergers, with appropriately adjusted energy and timescales: a He-burning shell will ignite off-center and diffuse inwards, possibly with episodic flashes due to the inefficiency of neutrino cooling \citep{iben90,sn98}." +. Once the burning shell reaches the center. the merger renmant appears as a core He-burnimg sdB or «dO star aud eventually cools to become a C/O WD.," Once the burning shell reaches the center, the merger remnant appears as a core He-burning sdB or sdO star and eventually cools to become a C/O WD." + The merger of a Πο and a C/O WD should vield a Ue shell buruig R CrB star (Webbink1981:Ibou1990:Clayton 1996).. although the final Hoe shell flash of cooling post-AGB stars las also been proposed as an R CrD progenitor (Ibenetal.19," The merger of a He and a C/O WD should yield a He shell burning R CrB star \citep{webb84,iben90,clay96}, although the final He shell flash of cooling post-AGB stars has also been proposed as an R CrB progenitor \citep{ity96}." +96).. θατοις of these uecreers within the framework of our updated physical uodel will allow us to connect theoretical preclictious with observational constraints (Asplundctal.1997:Claytonetal.2007.2011:Jeffery2011) aud help o differentiate between these progenitor scenarios.," Simulations of these mergers within the framework of our updated physical model will allow us to connect theoretical predictions with observational constraints \citep{aspl97,clay07,clay11,jks11} and help to differentiate between these progenitor scenarios." + We thank the anouvimous referee for comments that welped to guide and clarity our work., We thank the anonymous referee for comments that helped to guide and clarify our work. + We are very erateful o M. Dan aud S. Rosswoe for access to their simulation data., We are very grateful to M. Dan and S. Rosswog for access to their simulation data. + We thank them. R. Foley. W. IHaxton. D. Metzger. T. Piro. P. Podsiadlowski. J. Schliwab. S. Sim. aud. M. van Ierkwijk for helpful discussious.," We thank them, R. Foley, W. Haxton, B. Metzger, T. Piro, P. Podsiadlowski, J. Schwab, S. Sim, and M. van Kerkwijk for helpful discussions." + We also thank D. Paxton aud the MESA Council for providing iux irelesslw developing MESA., We also thank B. Paxton and the MESA Council for providing and tirelessly developing MESA. + This work was supporte w the National Science Foundation uuder grauts PITY 55-51161 and AST-11-0017L. the U.S. Department of Energy uncer erant DE-SC'000165I8. andthe Theoretica Astrophysics Center at UC Berkeley.," This work was supported by the National Science Foundation under grants PHY 05-51164 and AST-11-09174, the U.S. Department of Energy under grant DE-SC00046548, and the Theoretical Astrophysics Center at UC Berkeley." + KJS is supporte w NASA through the Eiusteiu Postdoctoral Fellowship awarded by the Chandra X-ray Center. which is operatcc x the Sinithsonian Astrophysical Observatory for NASA under contract NASS-03060.," KJS is supported by NASA through the Einstein Postdoctoral Fellowship awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060." + EQ is supported iu part by he David aud Lucile Packard Foundation., EQ is supported in part by the David and Lucile Packard Foundation. + DINs work is supported bv the Director. Office of Energy. Research. Office of Hlieh. Energy. aud Nuclear Physics. Divisions of Nuclear Physics. of the U.S. DOE under eraut. DE-ACU2-05CII11231 and by the DOE SciDAC Program (DE-FC'02-06ERLE138).," DK's work is supported by the Director, Office of Energy Research, Office of High Energy and Nuclear Physics, Divisions of Nuclear Physics, of the U.S. DOE under grant DE-AC02-05CH11231 and by the DOE SciDAC Program (DE-FC02-06ER41438)." + Iu this Appendix. we explain the computational details of our heieht-integrated viscous simulation whose results are shown in 83..," In this Appendix, we explain the computational details of our height-integrated viscous simulation whose results are shown in \ref{sec:viscevol}." + The approsimatious outlined there vield the Lagrangian time derivatives: DUC 2&jun )and( where the cvlindvically euclosed mass 15 n.," The approximations outlined there yield the Lagrangian time derivatives: = -, = v_r, = + ^2 - _K^2, - +, = - - 2 + 2 ^2 ^2, where the cylindrically enclosed mass is $m$." + The heieglt-inteerated quautitics X aud II are defined at zone ceuters. while the mass. radius. ορ aud ο are defined at zone bouudaries," The height-integrated quantities $\Sigma$ and $\Pi$ are defined at zone centers, while the mass, radius, $v_r$, and $\Omega$ are defined at zone boundaries." +" The radius at the inner boundary of the exid. where the euclosed mass is 0.9.AZ... is asstuned to be fixed: Ίνα, the degenerate core ds assunied to be incompressible."," The radius at the inner boundary of the grid, where the enclosed mass is $0.9 \msol$, is assumed to be fixed; i.e., the degenerate core is assumed to be incompressible." + The angular velocity eradicut across the immer aud outer bouudaries is set to zero., The angular velocity gradient across the inner and outer boundaries is set to zero. + These boundary couditious. while clearly approximations. do conserve the total cncrey and angular momentum in the erid to better than during the simulation.," These boundary conditions, while clearly approximations, do conserve the total energy and angular momentum in the grid to better than during the simulation." + We approxiuate the radial gravitational acceleration at (rat) Where + is the height from the midplane. as the value at the midplane.- frei.t)2—Gin4rz.," We approximate the radial gravitational acceleration at $(r_{\rm cyl},z)$, where $z$ is the height from the midplane, as the value at the midplane, $f_r(r_{\rm cyl}, z) \simeq -G m / r_{\rm cyl}^2 $." + Thisul is equivaleut to the assumption that Ixpid.10°. even using tree-based aleoritluus on the most powerful The situation is even worse for other N-body algorithuus.," Nevertheless, the computational expense has remained prohibitive for $ N > 10^9$, even using tree-based algorithms on the most powerful The situation is even worse for other N-body algorithms." + The N-body direct evolution cetloc scales as O(N?). which 1iakes it impossible to run simulations with more than Lot particles.," The N-body direct evolution method scales as $O(N^2)$, which makes it impossible to run simulations with more than $10^4$ particles." + To overcome this difficulty. and when high accuracy is required. alternative numerical iiethods basce on lierarchical foree-computation aleorithius are widely used.," To overcome this difficulty, and when high accuracy is required, alternative numerical methods based on hierarchical force-computation algorithms are widely used." + The recent effort has addressed the production of new software and algoritlins for the new ecucration of high-performance computer systems., The recent effort has addressed the production of new software and algorithms for the new generation of high-performance computer systems. + The ultimate target is an dplemeutation of the tree N-body algorithia to rui simulations with hieher accuracy aud particle uuuber. decreasing the cost of the simulation in terms of CPU time and increasing performance m terms of umber of particles/sccoud claborated when ruunuius on AIPP systems.," The ultimate target is an implementation of the tree N-body algorithm to run simulations with higher accuracy and particle number, decreasing the cost of the simulation in terms of CPU time and increasing performance in terms of number of particles/second elaborated when running on MPP systems." + Among the tree algoritaus designed to compute the gravitational force in N-body systems. oue of the most used aud powerful in modern cosinology is that by Barnes aud Thit (BID)|5].," Among the tree algorithms designed to compute the gravitational force in N-body systems, one of the most used and powerful in modern cosmology is that by Barnes and Hut (BH)." +. The DII octal-tree recursive inethod is inhereutly adaptive aud allows one to achieve a higher mass resolution even if parallel iuplementation of this algorithun suffers. from a serious drawback: it can casily run iuto inbalauce as soon as the coufiguratiou evolves. causing performance degradation.," The BH octal-tree recursive method is inherently adaptive and allows one to achieve a higher mass resolution even if parallel implementation of this algorithm suffers from a serious drawback: it can easily run into imbalance as soon as the configuration evolves, causing performance degradation." + In this paper we xeseut a modified version of the BIT aleorithim in which we have introduced an cubanced groupiug strateey., In this paper we present a modified version of the BH algorithm in which we have introduced an enhanced grouping strategy. + We will show how this feature allows an increase in performance when we consider Sinmlation with a large muuber of particles (N2» 10°).," We will show how this feature allows an increase in performance when we consider N-body simulation with a large number of particles $ N \geq +10^6$ )." + The code we preseut incorporates fully seriodic boundary conditions using the Ewald method. without the use of fast Fouricr transform echuiques Iu Section 2 we give a brief description of our N-body parallel code. based on the DIT tree algorithia. and the dynamic load balauce (DLD) policy adopted.," The code we present incorporates fully periodic boundary conditions using the Ewald method, without the use of fast Fourier transform techniques In Section 2 we give a brief description of our N-body parallel code, based on the BH tree algorithm, and the dynamic load balance (DLB) policy adopted." + Iu Section 3 we describe our euliauced erouping strategv., In Section 3 we describe our enhanced grouping strategy. + In Section Lowe show the results of our tests and in Section 5 we report our conclusions., In Section 4 we show the results of our tests and in Section 5 we report our conclusions. + Since the publication of the monograph by Hockuey and Eastwood in 1981 |17].. a new class of particle simulation methods," Since the publication of the monograph by Hockney and Eastwood in 1981 , a new class of particle simulation methods" +of interest 10) Myr),of interest $-100$ Myr). + Tn our best-fit model the major spiral arius of fie outer disk are located on the trailing side (lower pane of Fig. 2)))., In our best-fit model the major spiral arms of the outer disk are located on the trailing side (lower panel of Fig. \ref{fig:face-on}) )). + A model cube (right ascension. cdeclination. radial velocity) with the properties of the observations of Clune et al. (," A model cube (right ascension, declination, radial velocity) with the properties of the observations of Chung et al. (" +"2009) is produced frou, the model eas paricle distribution.",2009) is produced from the model gas particle distribution. + The observed aud inodel gas distribuious auc associated position velocity diagrams are show 1i Fig. 3.., The observed and model gas distributions and associated position velocity diagrams are shown in Fig. \ref{fig:ngc4330.pvd}. + The model gas distribution reproduces qualitatively the main features of he observations: {1) ie truncation of the ealactic eas ¢isk to the nortleast ud southwest. Gi) the gas distribuion along the minor xis is lavecr on the shadowed side {southeast). are (iii) re existence of the soutlivesteru ail. especially here is a mnodel counterpart (a kink) to the discoutiuuiv of le gas distribution between the disk aud the tail eus.," The model gas distribution reproduces qualitatively the main features of the observations: (i) the truncation of the galactic gas disk to the northeast and southwest, (ii) the gas distribution along the minor axis is larger on the shadowed side (southeast), and (iii) the existence of the southwestern tail, especially there is a model counterpart (a kink) to the discontinuity of the gas distribution between the disk and the tail gas." +ons Despite he overall agreement with oservatious. therο are disagreeueuts: 1) the upturn to the southeast is LOLC xonounced in the model compare to observations.," Despite the overall agreement with observations, there are disagreements: (i) the upturn to the southeast is more pronounced in the model compared to observations." + In articular. the low cohuun density eas on the eastern side of he galaxy is more extended than the observed eas distribution: (1) the structure of the πο) tail has a pronounced part which is bent by au auele of 18? with respect to the galactic disk. whereas the observed tail bends by an angle of ~30° (ui) the observed vertical extent of the disk is about twice as large as that of the model gas disk.," In particular, the low column density gas on the eastern side of the galaxy is more extended than the observed gas distribution; (ii) the structure of the model tail has a pronounced part which is bent by an angle of $18^{\circ}$ with respect to the galactic disk, whereas the observed tail bends by an angle of $\sim 30^{\circ}$; (iii) the observed vertical extent of the disk is about twice as large as that of the model gas disk." + The observed eas kinematics alowο the major axis are well reproduced by the model. excep that the model las au jutriusically stecver rotation curve and a somewhat Heh rotation velocity.," The observed gas kinematics along the major axis are well reproduced by the model, except that the model has an intrinsically steeper rotation curve and a somewhat high rotation velocity." + The exteusiois to high velocities observed in the position velocity «jaegranus of the eas ail are also present in the model., The extensions to high velocities observed in the position velocity diagrams of the gas tail are also present in the model. + However. there are differences between the observed iux the model position velocity diagrams aloug the ninor axis (4) the observed vosition velocity diagram along the nior axis at the ealaxy center and to the northeas show the highest inewidths of ~150 |: (ii) the‘corresponding model vosition velocity diagrais are two fines broader: (111) in he upturn regiong (northeast) the olServed velocities ou he windward side are. if they are different. hieher than hose of the opposite side. whereas he model velocities are lower.," However, there are differences between the observed and the model position velocity diagrams along the minor axis: (i) the observed position velocity diagram along the minor axis at the galaxy center and to the northeast show the highest linewidths of $\sim 150$ $^{-1}$; (ii) the corresponding model position velocity diagrams are two times broader; (iii) in the upturn region (northeast) the observed velocities on the windward side are, if they are different, higher than those of the opposite side, whereas the model velocities are lower." + Thus. whereas the moclel kineniaties of the tail," Thus, whereas the model kinematics of the tail" +obtained over a large area.,obtained over a large area. +" These data show a wide range of visual magnitudes for NGC 1333, and a non-Gaussian CO spectral profile consistent with multi-velocity components."," These data show a wide range of visual magnitudes for NGC 1333, and a non-Gaussian CO spectral profile consistent with multi-velocity components." +" These results are consistent and likely related to a layered cloud structure along the line of sight, which was first proposed by Ungerechts&Thaddeus(1987)."," These results are consistent and likely related to a layered cloud structure along the line of sight, which was first proposed by \citet{Unge87}." +. Interstellar extinction studies of field stars toward NGC 1333 also suggest at least two components in the line of sight at different distances toward NGC 1333 (Cernis1990)., Interstellar extinction studies of field stars toward NGC 1333 also suggest at least two components in the line of sight at different distances toward NGC 1333 \citep{Cernis90}. +". According to column density maps of the Perseus cloud (Ridgeetal.2006b),, the region studied here lies in the lower density envelope of NGC 1333."," According to column density maps of the Perseus cloud \citep{Ridge06b}, the region studied here lies in the lower density envelope of NGC 1333." +" Maps of high density molecular tracers (N;2H*, HCO*) as well as of the 870 jum dust emission, show that around IRAS 4A the dense gas has a filamentary distribution oriented in the NW-SE direction, with the long axis positioned at ~142°(Sandell&Knee2001;Olmietal.2005;Walsh2007)."," Maps of high density molecular tracers $_{2}$ $^{+}$, $^{+}$ ) as well as of the 870 $\mu$ m dust emission, show that around IRAS 4A the dense gas has a filamentary distribution oriented in the NW–SE direction, with the long axis positioned at $\simeq$ \citep{Sandell01, Olmi05, Walsh07}." +. The near-IR and optical polarization vectors of the background stars shown in Fig., The near-IR and optical polarization vectors of the background stars shown in Fig. +" 6 trace the POS component of the magnetic field associated with the lower density envelope around IRAS 4A/4B. South of these sources, where we have most of the polarization sample, the magnetic field has a direction of ~160°."," \ref{mapaVIS_IR} trace the POS component of the magnetic field associated with the lower density envelope around IRAS 4A/4B. South of these sources, where we have most of the polarization sample, the magnetic field has a direction of $\simeq 160\degr$." + The observed configuration is consistent with the results obtained at much larger scale by Goodmanetal.(1990) and Tamuraetal.(1988)., The observed configuration is consistent with the results obtained at much larger scale by \citet{Goodman90} and \citet{Tamura88}. +. According to the COMPLETE survey (Ridgeetal.2006a) the polarization was measured toward regions with a visual extinction of 4 to 5 mag., According to the COMPLETE survey \citep{Ridge06a} the polarization was measured toward regions with a visual extinction of 4 to 5 mag. + The magnetic field orientation derived from our data is roughly parallel to the dense filamentary structure associated with IRAS 4A (Walshetal.2007;Sandell&Knee2001)..," The magnetic field orientation derived from our data is roughly parallel to the dense filamentary structure associated with IRAS 4A \citep{Walsh07,Sandell01}." +" However, the submm polarization maps towards IRAS 4A and IRAS 4B show that the magnetic field within the filament is approximately perpendicular to the filament's major axis"," However, the submm polarization maps towards IRAS 4A and IRAS 4B show that the magnetic field within the filament is approximately perpendicular to the filament's major axis" +We note that scenario I — radio brightening which images reveal to be displaced from the binary core — has been directly observed in several cases (e.g. Corbel et al.,We note that scenario I – radio brightening which images reveal to be displaced from the binary core – has been directly observed in several cases (e.g. Corbel et al. + 2002: Gallo et al., 2002; Gallo et al. + 2004)., 2004). + In this scenario. the radio emission we observe in the soft state should in most cases have an optically thin radio spectrum. in contrast to the flat radio spectrum associated with the hard state (Fender 2001).," In this scenario, the radio emission we observe in the soft state should in most cases have an optically thin radio spectrum, in contrast to the flat radio spectrum associated with the hard state (Fender 2001)." + In Fig 6 we plot the radio spectral index for observations of H1743-32? as a function of X-ray hardness. which support a trend towards more negative spectral indices. indicative of a lower optical depth. with decreasing hardness.," In Fig \ref{spindex} we plot the radio spectral index for observations of H1743-322 as a function of X-ray hardness, which support a trend towards more negative spectral indices, indicative of a lower optical depth, with decreasing hardness." + In addition such emission should be observed to monotonically fade as the ejecta expand in the surrounding medium., In addition such emission should be observed to monotonically fade as the ejecta expand in the surrounding medium. + Occasional rebrightenings and spectral flattenings can be associated with shocks. whether internal or external.," Occasional rebrightenings and spectral flattenings can be associated with shocks, whether internal or external." + Scenario II — the existence. sometimes. of a powerful jet from the core in the soft state — cannot be ruled out. but seems to be the more complex solution. requiring explanations for why the je spectrum has switched to optically thin. and why it only happens sometimes.," Scenario II – the existence, sometimes, of a powerful jet from the core in the soft state – cannot be ruled out, but seems to be the more complex solution, requiring explanations for why the jet spectrum has switched to optically thin, and why it only happens sometimes." + We conclude that the scenario I. which agrees with the curren yaradigm of a suppressed (quenched) core jet in the soft state is more likely. although scenario IT cannot be ruled out.," We conclude that the scenario I, which agrees with the current paradigm of a suppressed (`quenched') core jet in the soft state is more likely, although scenario II cannot be ruled out." + In order to tes his further. sensitive high resolution (VLBI) and/or high frequency (mm/IR) observations are required throughout an outburst.," In order to test this further, sensitive high resolution (VLBI) and/or high frequency (mm/IR) observations are required throughout an outburst." + Note hat the near-IR evidence for suppression of the jet in the sof state (Homan et al., Note that the near-IR evidence for suppression of the jet in the soft state (Homan et al. + 2005. Russell et al.," 2005, Russell et al." + 2007) clearly supports he case that the flat-spectrum jet is suppressed. and therefore tha he jet power is significantly reduced (whether the jet emission is suppressed at all frequencies or just becomes optically thin).," 2007) clearly supports the case that the flat-spectrum jet is suppressed, and therefore that the jet power is significantly reduced (whether the jet emission is suppressed at all frequencies or just becomes optically thin)." + It is clear from many radio observations (e.g. Tanabaum et al., It is clear from many radio observations (e.g. Tanabaum et al. + 1974: Fender et al., 1974; Fender et al. + 1999: Corbel et al., 1999; Corbel et al. + 2000. 2001: Gallo. Fender Pooley 2003) that the radio emission from BHXRBs reactivates when the sources transition back to the hard state following a period in a softer state.," 2000, 2001; Gallo, Fender Pooley 2003) that the radio emission from BHXRBs reactivates when the sources transition back to the hard state following a period in a softer state." +" In FBGOA it was suggested (speculatively) that there may be a vertical ""jet line’ for each HID (and maybe each source).", In FBG04 it was suggested (speculatively) that there may be a vertical `jet line' for each HID (and maybe each source). + In harder states than this jet line a steady jet would be produced. and in softer states the jet suppressed.," In harder states than this jet line a steady jet would be produced, and in softer states the jet suppressed." + Transition across the jet line from harder to softer states (i.e. right to left) would result in a radio flare (due to internal shocks as jet velocity increased). whereas the transition from soft to hard (left to right: at lower luminosity) would simply produce the reactivation of the (flat spectrum) jet without a large flare.," Transition across the jet line from harder to softer states (i.e. right to left) would result in a radio flare (due to internal shocks as jet velocity increased), whereas the transition from soft to hard (left to right; at lower luminosity) would simply produce the reactivation of the (flat spectrum) jet without a large flare." + It seems clear that the jet does remain on until at least the moment of the radio flare. i.e. in the hard-intermediate state. during the overall hard + soft transition (see Fig |: also FBGOA and Corbel et al.," It seems clear that the jet does remain on until at least the moment of the radio flare, i.e. in the hard-intermediate state, during the overall hard $\rightarrow$ soft transition (see Fig 1; also FBG04 and Corbel et al." + 2004)., 2004). + However. data testing the reactivation of the hard state steady jet during the soft to hard transition has been hard to come by.," However, data testing the reactivation of the hard state steady jet during the soft to hard transition has been hard to come by." + Comparison of the 1996-1999 and 2002-2003 outbursts of GX 339-4 indicates that upper limits to the jet flux have been measured. in 1909. at almost exactly the same hardness ratio where the mayor radio flare was observed in 2002.," Comparison of the 1996–1999 and 2002–2003 outbursts of GX 339-4 indicates that upper limits to the jet flux have been measured, in 1999, at almost exactly the same hardness ratio where the major radio flare was observed in 2002." + This is shown in some detail in Fig7., This is shown in some detail in Fig. + Note that there may be some uncertainties associated with comparing colours at different epochs due to changes in the response of the RATE PCAs., Note that there may be some uncertainties associated with comparing colours at different epochs due to changes in the response of the RXTE PCAs. + The upper limits to the radio flux density in 1999 are clearly also. however. at much lower X-ray count rates.," The upper limits to the radio flux density in 1999 are clearly also, however, at much lower X-ray count rates." + Could this be an explanation for the non-detection?, Could this be an explanation for the non-detection? + The upper limits of <0.2 mJy were obtained at a count rate of ~SO ct/sec., The upper limits of $\leq 0.2$ mJy were obtained at a count rate of $\sim 80$ ct/sec. + The ~12 mly steady detection and subsequent flare were observed at a much higher count rate. of ~1300 ct/sec.," The $\sim 12$ mJy steady detection and subsequent flare were observed at a much higher count rate, of $\sim 1300$ ct/sec." +" Combining these two results. a relation between radio and X-ray fluxes steeper than Pii,xpl would be required to account for the non-detection in 1999."," Combining these two results, a relation between radio and X-ray fluxes steeper than $F_{\rm radio} \propto +F_{\rm X}^{1.4}$ would be required to account for the non-detection in 1999." + This is of course far steeper than the fiisxUS relation of Corbel et al. (," This is of course far steeper than the $F_{\rm radio} \propto +F_{\rm X}^{\sim 0.7}$ relation of Corbel et al. (" +2002) for this and other (Gallo et al.,2002) for this and other (Gallo et al. + 2003) sources in the hard state., 2003) sources in the hard state. + ΑΕΕ index is in fact as expected for jets from a radiatively efficient accretion flow (e.g. Heinz Sunyaev 2003: see Migliari Fender 2006 for the possible case of such an index for neutron stars) and so this is not quite ruled out., A +1.4 index is in fact as expected for jets from a radiatively efficient accretion flow (e.g. Heinz Sunyaev 2003; see Migliari Fender 2006 for the possible case of such an index for neutron stars) and so this is not quite ruled out. + Nevertheless. it seems likely that this combination of observations implies that a vertical “jet line’. implying a one-to-one correspondence between X-ray hardness and core jet properties. does not exist.," Nevertheless, it seems likely that this combination of observations implies that a vertical `jet line', implying a one-to-one correspondence between X-ray hardness and core jet properties, does not exist." + Of course there is nopriori reason to assume a vertical jet line. but since this is what was sketched in FBGOA we need to address this point.," Of course there is no reason to assume a vertical jet line, but since this is what was sketched in FBG04 we need to address this point." + As noted in Corbel et al. (, As noted in Corbel et al. ( +2000). the first redetection of GX,"2000), the first redetection of GX" +euvelopes. except that now the We euvelopes from the two stars meet aud are expelled.,"envelopes, except that now the He envelopes from the two stars meet and are expelled." + For other possible evolutionary scenarios of NS-NS binaries. please refer Yamaokaetal.(1993):Fryver&Kalogera(1997):ArzoumanianFraucischellietal.(2002):Dewi&Pols(2003):Willems(200D:Thorsett 2005].," For other possible evolutionary scenarios of NS-NS binaries, please refer \citet{Yam93,Fry97,Arz99,Fra02,Dew03,Wil04,Tho05}." + Using the four acemvately measured companion masses in relativistic binary NS«4 we coustruct the final pulsar masses as in Fie., Using the four accurately measured companion masses in relativistic binary NS's we construct the final pulsar masses as in Fig. + 2 and Table dL., \ref{fig2} and Table \ref{tab2}. + Our strateey is to calculate the mass acereted on the primary pulsar in the Π τος eiut stage of secondary star., Our strategy is to calculate the mass accreted on the primary pulsar in the H red giant stage of secondary star. + There cau be an extra accretion in the We red giant stage. which will increase the final primary compact star mass.," There can be an extra accretion in the He red giant stage, which will increase the final primary compact star mass." + We also neglected mass transter by We winds in the pulsar. We star binary which precedes the NS-NS stage. but this is likely to be small because of the propeller effect (Francischellietal.2002).," We also neglected mass transfer by He winds in the pulsar, He star binary which precedes the NS-NS stage, but this is likely to be small because of the propeller effect \citep{Fra02}." +. For the relation between the pulsar and progenitor masses. we assumed Note that this approximation is based ou the known pulsar masses. as in Table 2.. and the pulsar and progenitor mass relations. which we discussed above aud marked in Table 1..," For the relation between the pulsar and progenitor masses, we assumed Note that this approximation is based on the known pulsar masses, as in Table \ref{tab3a}, and the pulsar and progenitor mass relations, which we discussed above and marked in Table \ref{tab2}. ." + Usine the Bethe&Brown(1998) hvpercritical accretion corrected by Belezvnskietal.(2002) we arrive at Fig.," Using the \citet{Bet98} hypercritical accretion corrected by \citet{Bel02} + we arrive at Fig." + and Table 1.. ceiving the estimated fal primary compact star masses.," \ref{fig2} and Table \ref{tab2}, giving the estimated final primary compact star masses." + Note that there can be extra mass accretion during Ile eiut stage. OLAS.0.2M.. therefore. final compact star masses iu Fie.," Note that there can be extra mass accretion during He giant stage, $0.1 \msun - 0.2\msun$, therefore, final compact star masses in Fig." + 2 aud Table 1. can be treated as a lower linüt., \ref{fig2} and Table \ref{tab2} can be treated as a lower limit. + Iu Table 2 we list the known relativistic NS-NS binaries., In Table \ref{tab3a} we list the known relativistic NS-NS binaries. + The masses of the two least massive binaries 3039A. D aud 2251 are slightly differeut because of mass transter during the Te red giant stage.but certainly the near," The masses of the two least massive binaries $-$ 3039A, B and $-$ 2251 are slightly different because of mass transfer during the He red giant stage,but certainly the near" +where Jy and {ο are the outer and (he inner equatorial axis length. A is (he ratio of the length of the equatorial axis to the polar axis. 0; is the viewing augle of (he observer. and ji=cos 8.,"where $R_1$ and $R_2$ are the outer and the inner equatorial axis length, $A$ is the ratio of the length of the equatorial axis to the polar axis, $\theta_i$ is the viewing angle of the observer, and $\mu=\cos\theta$ ." +" At an edge-on view. 0;=7/2 and ó=0. and hence Nj, is real."," At an edge-on view, $\theta_i=\pi/2$ and $\phi=0$, and hence $N_{lm}$ is real." + From the equation of hydrostatic equilibrium. we have an expression of the particle densitv as a function of gas pressure: where P? is the pressure at different geometrical height. pis (he mass density at different pressure scale. and g is the surface gravity (which can be assumed to be constant [or a eeometrically (hin atmosphere).," From the equation of hydrostatic equilibrium, we have an expression of the particle density as a function of gas pressure: where $P$ is the pressure at different geometrical height, $\rho$ is the mass density at different pressure scale, and $g$ is the surface gravity (which can be assumed to be constant for a geometrically thin atmosphere)." + Substituting eqs., Substituting eqs. + 11. aud 12. into eq. 10..," \ref{abc2a} and \ref{abc3} into eq. \ref{abc2}," + we have for the degree of polarization lor a hydrostatic atimosphliere viewed edge on., we have for the degree of polarization for a hydrostatic atmosphere viewed edge on. + The effective temperature of the L dwarts of different spectral (wpe can be approximated bv the linear empirical relationship of Stephensetal.(2001): where L is the spectral (wpe between LO ancl Ls., The effective temperature of the L dwarfs of different spectral type can be approximated by the linear empirical relationship of \citet{ste01}: where $L$ is the spectral type between L0 and L8. + Recently Vrbaetal.(2004) and Golimowskietal.(2004) have presented {ει measurements for L and T dwarls based on bolometric Iuminosities., Recently \citet{vrba04} and \citet{goli04} have presented $T_{eff}$ measurements for L and T dwarfs based on bolometric luminosities. +" Ii the present work we adopt the sixth order polynomial fit given by Golimowskiοἱal.(2004). for translating the Tip, into spectral type.", In the present work we adopt the sixth order polynomial fit given by \citet{goli04} for translating the $T_{eff}$ into spectral type. + Fig., Fig. + 1 shows the Τε for different spectral (vpes calculated by using the linear relationship given by Stephens and by the polynomial formula provided by Golimowskietal...(2004)., 1 shows the $T_{eff}$ for different spectral types calculated by using the linear relationship given by \citet{ste01} and by the polynomial formula provided by \citet{goli04}. +. The T;jy calibration of Vrbaetal.(2004). and Golimowskietal.(2004) agree well in the interval L3-L8 but there are significant differences in earlier (vpes., The $T_{eff}$ calibration of \citet{vrba04} and \citet{goli04} agree well in the interval L3-L8 but there are significant differences in earlier types. +" In our caleulations for (he degree of polarization. the effectivetemperature 7;;, is used and hence the degree of polarization should be considered strictly as a function of 7;;; rather (han the spectral (vpe."," In our calculations for the degree of polarization, the effectivetemperature $T_{eff}$ is used and hence the degree of polarization should be considered strictly as a function of $T_{eff}$ rather than the spectral type." +Tosi 1996).,Tosi 1996). + We have computed a model for the solar neighborhood uuder the same assunuptious as those adopted by Cariei (1996). but based on the vields by WLW and WW. aud adoptedeam a 0.01 8$ , enrich the ISM with oxygen." + e) For SNIa we have aken iuto account the vieldsby Nomoto. Thiclemamn. Yokoi (198L. model W7).," g) For SNIa we have taken into account the yieldsby Nomoto, Thielemann, Yokoi (1984, model W7)." + Oulv a fraction of binary stars. m the 3?πι|oma<16 ranec. become SNIDI where 2.2$ 15 Gyr) only their line-strengths are consistent $\sigma$ uncertainty) with this oldest isochrone (solid line in Figure \ref{fig:old.globulars}) ). + Otherwise. SWB VIL clusters with strong Balmer-absorption are given vounger spectroscopic ages.," Otherwise, SWB VII clusters with strong Balmer-absorption are given younger 'spectroscopic' ages." + The age predictions of the SSP models for the LMC elobular clusters (Figure 15)) show significant variations. depending upon the combination of Lick/LDS indices used.," The age predictions of the SSP models for the LMC globular clusters (Figure \ref{fig:age}) ) show significant variations, depending upon the combination of Lick/IDS indices used." + ln contrast. the literature ages of the clusters (5 from the CCMDs of (1998). 1 from the eround-based CAID of (1997) and 2 inferred. from their SWB type) show a reasonably tight age-range. from 10 to 17 Gyr.," In contrast, the literature ages of the clusters (5 from the CMDs of (1998), 1 from the ground-based CMD of (1997) and 2 inferred from their SWB type) show a reasonably tight age-range, from 10 to 17 Gyr." + These ages are in reasonable agreement with the “best” value for the Milky Way GC's of 12.0 + 2.9 Gyr(7)., These ages are in reasonable agreement with the 'best' value for the Milky Way GCs of 12.9 $\pm$ 2.9 Gyr. +. We find that the higher-order Balmer lines. σι ancl 11δι-. are generally consistent with the CALD turn-oll ages and/or integrated colours.," We find that the higher-order Balmer lines, $\gamma_{\rm F}$ and $\delta_{\rm F}$, are generally consistent with the CMD turn-off ages and/or integrated colours." + However. 117. the most. age-sensitive of the Lick/LDS indices predicts a couple of clusters to be significantly than is indicated by their literature values.," However, $\beta$, the most age-sensitive of the Lick/IDS indices predicts a couple of clusters to be significantly than is indicated by their literature values." + Specifically. the two clusters. NGC 1754 and NGC 2005. have SSP model ages which are inconsistcnt with the literature at 30 significance.," Specifically, the two clusters, NGC 1754 and NGC 2005, have SSP model ages which are inconsistent with the literature at $\sigma$ significance." + Both of these clusters have well-developect blue LBs. and their influence upon the integrated. Balmer indices we examine shortly.," Both of these clusters have well-developed blue HBs, and their influence upon the integrated Balmer indices we examine shortly." + As shown in Figure 16.. we also find that the position of NGC 1786 stands out in the σι and Loy planes of the SSP models.," As shown in Figure \ref{fig:old.globulars}, we also find that the position of NGC 1786 stands out in the $\gamma_{\rm F}$ and $\delta_{\rm F}$ planes of the SSP models." + However. these indices of this cluster only fall below the model isochrones at the ~ 26 level. and therefore we do not aseribe them any particular significance.," However, these indices of this cluster only fall below the model isochrones at the $\sim$ $\sigma$ level, and therefore we do not ascribe them any particular significance." + In the previous section. we found two GCs. NGC 1754 and NGC 2005. have SSP-derived ages significantly (30) vounger ian the literature values. (," In the previous section, we found two GCs, NGC 1754 and NGC 2005, have SSP-derived ages significantly $\sigma$ ) younger than the literature values. (" +1982) was the first to suggest wt the presence of blue LB stars may have a significant cect upon the equivalent width. of the Balmer indices. possibly comparable to the contribution from stars at the main sequence turn-olf.,"1982) was the first to suggest that the presence of blue HB stars may have a significant effect upon the equivalent width of the Balmer indices, possibly comparable to the contribution from stars at the main sequence turn-off." + Since NGC 2005 has a largely blue LIB. is it. possible that the presence of blue LLB stars are the origin of this present. disagreement with the SSP models?," Since NGC 2005 has a largely blue HB, is it possible that the presence of blue HB stars are the origin of this present disagreement with the SSP models?" + Looking at this issue. (1995) compared. the LL? Gine-strengths of 10 Galactic GCs with their LLB morphologies.," Looking at this issue, (1995) compared the $\beta$ line-strengths of 10 Galactic GCs with their HB morphologies." + With the aicl of empirical modelling. these authors suggested wt Blue LBs may inerease H2 by upwards of ~ 1.0A.," With the aid of empirical modelling, these authors suggested that Blue HBs may increase $\beta$ by upwards of $\sim$ 1.0." +. For the five clusters in our sample with ‘CAIDs from (1998). we have an accurate measure of wir CAID age. LIB morphology and integrated Balmer indices.," For the five clusters in our sample with CMDs from (1998), we have an accurate measure of their CMD age, HB morphology and integrated Balmer indices." + Since these clusters (NGC L754. NGC 1835. NGC 1898. NGC 2005 and NGC 2019) all have similar mean spectroscopic metallicitics (mean metallicities derived from Table Bl are Fe/H] = L38. 1.57. 1.30. 143 and 1.42 respectively). we are able to directly compare the effect. of UB morphology upon the Lick/LDS Balmer indices of the LAIC clusters.," Since these clusters (NGC 1754, NGC 1835, NGC 1898, NGC 2005 and NGC 2019) all have similar mean spectroscopic metallicities (mean metallicities derived from Table B1 are [Fe/H] = –1.38, –1.57, –1.30, –1.43 and –1.42 respectively), we are able to directly compare the effect of HB morphology upon the Lick/IDS Balmer indices of the LMC clusters." + We plot the Lick/LDS Balmer indices of these GCs against their respective LB parameters in Figure 17.., We plot the Lick/IDS Balmer indices of these GCs against their respective HB parameters in Figure \ref{fig:hb}. + The LB parameter 1994) is commonly. given by (DHOPESVoy RA). where D ds the number of HD. stars to the blue of the instability strip. V the number of Itt Lyrae stars ancl Z2 the number of LIB stars recward of the instability strip.," The HB parameter 1994) is commonly given by $(B-R)/(B+V+R)$ , where $B$ is the number of HB stars to the blue of the instability strip, $V$ the number of RR Lyrae stars and $R$ the number of HB stars redward of the instability strip." + Negative values of this parameter indicate a red LB morphology (1 corresponds to no blue LLB stars). increasingly positive numbers correspond. to. progressively rluer LBs.," Negative values of this parameter indicate a red HB morphology (–1 corresponds to no blue HB stars), increasingly positive numbers correspond to progressively bluer HBs." + As shown in Figure 17.. there is a correlation between UB parameter and Balmer line-strength. in the sense that Xxuer LIBs Lead to larger index values for L2. Lisp and Lop.," As shown in Figure \ref{fig:hb}, there is a correlation between HB parameter and Balmer line-strength, in the sense that bluer HBs lead to larger index values for $\beta$, $\gamma_{\rm F}$ and $\delta_{\rm F}$ ." + GC. 1505. which has a roughly. equal number of stars to he left and right of its RR. Lyrac variables (ancl has the recldest determined LEB morphology in our sample). has H3 inc-strengths ~ 1 lower than that of NGC 2005. a cluster which has a blue IIB.," NGC 1898, which has a roughly equal number of stars to the left and right of its RR Lyrae variables (and has the reddest determined HB morphology in our sample), has $\beta$ line-strengths $\sim$ 1 lower than that of NGC 2005, a cluster which has a blue HB." + If we temporarily ignore the possible contribution from. LIB stars. at à given metallicity. a change in ~ | lin LL? would be interpreted: as an age dilference of 2 10 Car using the (extrapolated) W94 models.," If we temporarily ignore the possible contribution from HB stars, at a given metallicity, a change in $\sim$ 1 in $\beta$ would be interpreted as an age difference of $>$ 10 Gyr using the (extrapolated) W94 models." + Since (1998) obtained a CALD age for NCC 1808 of 14.0 4 2.3 Cove. and for NGC 2005 of 16.6 £ 5.1 Gyr (assuming 1991 abundances) the possibility ofa 10 Civr age dillerence between these two ‘Lusters seems unlikely.," Since (1998) obtained a CMD age for NGC 1898 of 14.0 $\pm$ 2.3 Gyr, and for NGC 2005 of 16.6 $\pm$ 5.1 Gyr (assuming 1991 abundances), the possibility of a 10 Gyr age difference between these two clusters seems unlikely." + One possible origin of this behaviour of the Balmer lines isthe metallicity dependence of the integrated. indices rermselyves., One possible origin of this behaviour of the Balmer lines isthe metallicity dependence of the integrated indices themselves. + We found in this study that the metallicities of, We found in this study that the metallicities of +afterglow. much steeper than that measured during the prompt event.,"afterglow, much steeper than that measured during the prompt event." + If we assume a similar. or even steeper. spectral slope in the 2-700 keV energy interval (uncorrected for redshift). the total energy output from the afterglow is substantial.," If we assume a similar, or even steeper, spectral slope in the 2-700 keV energy interval (uncorrected for redshift), the total energy output from the afterglow is substantial." + For the same power law index measured in the 2-10 keV energy band. we obtain a fluence in the afterglow 1.5«10©ergcnm57.," For the same power law index measured in the 2-10 keV energy band, we obtain a fluence in the afterglow $~4.5\times 10^{-6} \flux$." + Using. a more conservative assumption. adding an exponential cutoff with folding energy 7-50 keV. the total fluence in the afterglow up to 2.5 days after the main event is 10©erecni2.," Using a more conservative assumption, adding an exponential cutoff with folding energy $\sim$ 50 keV, the total fluence in the afterglow up to 2.5 days after the main event is $\times 10^{-6} \flux$." + We conclude that the X and 7 total luminosity in the afterglow may be estimated to be between «107 ergs and 107 ergs.," We conclude that the X and $\gamma$ total luminosity in the afterglow may be estimated to be between $\times +10^{52}$ ergs and $\times 10^{53}$ ergs." + This is comparable to the total radiated power in the main GRB event., This is comparable to the total radiated power in the main GRB event. + This estimate is only speculative. as the spectral data of the afterglow are consistent with any cutoff energy above ~2 keV confidence) but it points out the need for prompt measurements of the afterglow that can give a good estimate of the spectral shape above 10 Unfortunately the spectral evolution in hard X-rays (above 10 keV) of most GRBs seems unaccessible to the present generation of telescopes operating in this energy band.," This estimate is only speculative, as the spectral data of the afterglow are consistent with any cutoff energy above $\sim$ 2 keV confidence) but it points out the need for prompt measurements of the afterglow that can give a good estimate of the spectral shape above 10 Unfortunately the spectral evolution in hard X-rays (above 10 keV) of most GRBs seems unaccessible to the present generation of telescopes operating in this energy band." + This is a very important observational point that cannot be fulfilled up to the next generation of focusing hard X-ray telescopes. maybe at least for a decade.," This is a very important observational point that cannot be fulfilled up to the next generation of focusing hard X-ray telescopes, maybe at least for a decade." + The time resolved spectral analysis of the afterglow reported in Table |. even if of low statistical quality. do suggest that the spectral shape remains stable during our observation of the afterglow.," The time resolved spectral analysis of the afterglow reported in Table 1, even if of low statistical quality, do suggest that the spectral shape remains stable during our observation of the afterglow." + While these data cannot be profitably used to constrain the spectral evolution of the afterglow in X-rays. they confirm the stability of the that is producing the observed X-ray emission in spite of the substantial power-law decline with time in the observed X-ray luminosity.," While these data cannot be profitably used to constrain the spectral evolution of the afterglow in X–rays, they confirm the stability of the that is producing the observed X–ray emission in spite of the substantial power–law decline with time in the observed X–ray luminosity." + The spectral evolution from GRB to the X-ray afterglow indicates that the energy distribution shifts towards lower energies. as expected from theoretical models.," The spectral evolution from GRB to the X-ray afterglow indicates that the energy distribution shifts towards lower energies, as expected from theoretical models." + The model for spectral evolution of Sari et al. (1998)), The model for spectral evolution of Sari et al. \cite{sari}) ) + suggests that the high energy tail of the emitted photon spectrum of the afterglow has a power law slope ~r] in the case of fast cooling. where p is the index of a power law distribution of the electrons.," suggests that the high energy tail of the emitted photon spectrum of the afterglow has a power law slope $\sim-\frac{\rm p}{2} - 1$ in the case of fast cooling, where p is the index of a power law distribution of the electrons." + Our measurement of a power law slope ~—1.6 implies p~1.2. à value that would give a non-finite energy in the electrons.," Our measurement of a power law slope $\sim -1.6$ implies $\sim 1.2$, a value that would give a non-finite energy in the electrons." + In this case the observations can be reconciled with theory assuming that a suitable cutoff. e.g. exponential. in the energy distribution of the parent electron population is present.," In this case the observations can be reconciled with theory assuming that a suitable cutoff, e.g. exponential, in the energy distribution of the parent electron population is present." + The observation of this cutoff in the produced photon distribution is however beyond the capability of the present generation of X—ray telescopes and may be accessible to the new missions like and if it is substantially below 100 keV. Our data support (as suggested also by Waxman 1997a. 1997b)) that we observe a regime where the synchrotron cooling time is long compared to the dynamical time.," The observation of this cutoff in the produced photon distribution is however beyond the capability of the present generation of X--ray telescopes and may be accessible to the new missions like and if it is substantially below 100 keV. Our data support (as suggested also by Waxman \cite{waxman_1,waxman_2}) ) that we observe a regime where the synchrotron cooling time is long compared to the dynamical time." + In this regime TDO and for a finite power in the electrons one obtains DL«1.5. definitely compatible with our measurement.," In this regime $\Gamma = -\frac{({\rm p}-1)}{2} - 1$ , and for a finite power in the electrons one obtains $\Gamma < -1.5$, definitely compatible with our measurement." + A caveat must be added to this interpretation: as the adiabatic losses become dominant in this regime. the observed luminosity i$. produced via synchrotron losses in much an inefficient way.," A caveat must be added to this interpretation: as the adiabatic losses become dominant in this regime, the observed luminosity is produced via synchrotron losses in much an inefficient way." + This brings down the efficiency of the radiative process and in parallel rises accordingly the total energy budget in the afterglow., This brings down the efficiency of the radiative process and in parallel rises accordingly the total energy budget in the afterglow. + The slope of the NIR-to-X-ray afterglow spectrum. corrected for Galactic and local extinction. and its temporal evolution are in fair agreement with models of expanding fireballs.," The slope of the NIR-to-X-ray afterglow spectrum, corrected for Galactic and local extinction, and its temporal evolution are in fair agreement with models of expanding fireballs." + Our assumption on the intrinsic extinction reasonably conforms with the proposal that GRBs are connected with star formation. and therefore expected to reside in star forming regions of their host galaxies (though not necessarily in extreme starburst galaxies. see Odewahn et al. 1998)).," Our assumption on the intrinsic extinction reasonably conforms with the proposal that GRBs are connected with star formation, and therefore expected to reside in star forming regions of their host galaxies (though not necessarily in extreme starburst galaxies, see Odewahn et al. \cite{odewahn}) )." + We note that the local extinction correction we adopted has the advantage. with respect to other approaches. of using a specific model curve and of making the spectrum consistent with a single radiation component over more than three decades of frequency.," We note that the local extinction correction we adopted has the advantage, with respect to other approaches, of using a specific model curve and of making the spectrum consistent with a single radiation component over more than three decades of frequency." + If the emission is isotropic the total X-~ luminosity in GRB971214. if at redshift z=3.42. is quite higher than [001 ergs. as already pointed out by Kulkarni et al. (1998)).," If the emission is isotropic the total $\gamma$ luminosity in GRB971214, if at redshift z=3.42, is quite higher than $^{51}$ ergs, as already pointed out by Kulkarni et al. \cite{kulkarni_n}) )." + Our analysis shows that a substantial fraction. more than60%.. of the total radiated energy in 2-10 keV is in the afterglow.," Our analysis shows that a substantial fraction, more than, of the total radiated energy in 2-10 keV is in the afterglow." + In addition a reasonable guess of a high energy exponential cutoff. with a folding energy of 50 keV. brings us to conclude that the total power in GRB971214 may be grossly underestimated 1f based only on the prompt event.," In addition a reasonable guess of a high energy exponential cutoff, with a folding energy of 50 keV, brings us to conclude that the total power in GRB971214 may be grossly underestimated if based only on the prompt event." + If the efficiency to convert the total energy output from the GRB in the afterglow is low. e.g. ~10%.. as usually assumed in most theoretical models of fireball expansion. our measure of a substantial energy output in 2-10 keV during afterglow shows that the total energy balance of a GRB is grossly underestimated.," If the efficiency to convert the total energy output from the GRB in the afterglow is low, e.g. $\sim$, as usually assumed in most theoretical models of fireball expansion, our measure of a substantial energy output in 2–10 keV during afterglow shows that the total energy balance of a GRB is grossly underestimated." + Furthermore. if we consider the probable presence of a high energy tail of the afterglow. a conservative estimate may bring the total energy output from GRB971214 to more than 10°! erg for an isotropically emitting source.," Furthermore, if we consider the probable presence of a high energy tail of the afterglow, a conservative estimate may bring the total energy output from GRB971214 to more than $^{54}$ erg for an isotropically emitting source." +" Alternatively. one may more comfortably stay with a luminosity of 10? erg assuming a more efficient mechanism to effectively extract radiative power from the expanding fireball or assuming an extreme cutoff to the high energy spectrum,"," Alternatively, one may more comfortably stay with a luminosity of $^{53}$ erg assuming a more efficient mechanism to effectively extract radiative power from the expanding fireball or assuming an extreme cutoff to the high energy spectrum." + A way out of this deadlock may be beaming (Yi 190. Shaviv Dar 1995)).," A way out of this deadlock may be beaming (Yi \cite{yi}, , Shaviv Dar \cite{shdar}) )." + Different authors have discussed the “beaming solution” to the GRB/afterglow observational problem (Dar 1998.1999.. Rhoads 1997.. Mésszárros Rees 1997b.. Drodzova Panchenko 1997.. Panaitescu et al. 1998)).," Different authors have discussed the “beaming solution” to the GRB/afterglow observational problem (Dar \cite{dar1,dar2}, Rhoads \cite{rhoads}, Mésszárros Rees \cite{meszrees}, Drodzova Panchenko \cite{drodzova}, Panaitescu et al. \cite{panait}) )." + If the emitted power is strongly beamed. and therefore not isotropically distributed. the total power in the GRB may be reduced by orders of magnitude. depending on the beam open angle.," If the emitted power is strongly beamed, and therefore not isotropically distributed, the total power in the GRB may be reduced by orders of magnitude, depending on the beam open angle." + Of course this has some major and obvious impacts., Of course this has some major and obvious impacts. + The number of GRBs (not detected at earth) rises by the same orders of magnitude., The number of GRBs (not detected at earth) rises by the same orders of magnitude. + Limb darkening (due to the random distribution of viewing angles inside the emission cone). time dependent (due to different beaming at different times after the main event)and energy dependent (due to the time-dependent photon energy distribution) effects should be observable once the afterglow sample ts large enough.," Limb darkening (due to the random distribution of viewing angles inside the emission cone), time dependent (due to different beaming at different times after the main event)and energy dependent (due to the time–dependent photon energy distribution) effects should be observable once the afterglow sample is large enough." + Examplesof such effects are discussed, Examplesof such effects are discussed +Found. of this type of obscured object in dwarf. spheroidal galaxies.,found of this type of obscured object in dwarf spheroidal galaxies. + The nearest comparable objects. though they. are somewhat bluer((./.—A)~ 2.4). are the two carbon Miras that have been found in the Sagittarius dwarf spheroical (Whitelock et al.," The nearest comparable objects, though they are somewhat $(J-K) \sim +2.4$ ), are the two carbon Miras that have been found in the Sagittarius dwarf spheroidal (Whitelock et al." + 1999)., 1999). + A comparison of the data from the two epochs of observation (JD 2P451926.52 and 2452263.55) shows that a number of the brighter stars are variable., A comparison of the data from the two epochs of observation (JD 2451926.52 and 2452263.55) shows that a number of the brighter stars are variable. + We list these stars in Table 2. where we give the magnitude difference between the second and first epochs. together with the standard. deviation. of the cillerence based on the internal errors of measurement," We list these stars in Table 2, where we give the magnitude difference between the second and first epochs, together with the standard deviation of the difference based on the internal errors of measurement" +the physical parameters that exists in such a steady state.,the physical parameters that exists in such a steady state. +" We can express [ in terms of γε and v, using the pair creation threshold criteria (Eq. 3))", We can express $\G$ in terms of $\gamma_e$ and $\nu_\gamma$ using the pair creation threshold criteria (Eq. \ref{fKN}) ) +" and we can express vy, in terms of v, and γε (using Eq. 1)).", and we can express $\nu_L$ in terms of $\nu_\g$ and $\g_e$ (using Eq. \ref{gamma}) ). + Given these expressions we can estimate the steady state Yq as a function of γε., Given these expressions we can estimate the steady state $Y_H$ as a function of $\g_e$ . + Fig., Fig. + 5 depicts the resulting Y; values as a function of γε for different values of a., \ref{fig:Pairs} depicts the resulting $Y_H$ values as a function of $\g_e$ for different values of $\alpha$. +" The UV solution for vr,>10v;,; and with rather low values of γε and Τ is possible.", The UV solution for $\nu_L > 10 \nuopt$ and with rather low values of $\g_e$ and $\G$ is possible. + However this solutions suffers from the problems discussed earlier., However this solutions suffers from the problems discussed earlier. + It seems that if we impose the pair creation threshold conditions the IR solution is ruled out with very high Y; values (for any reasonable a)., It seems that if we impose the pair creation threshold conditions the IR solution is ruled out with very high $Y_H$ values (for any reasonable $\alpha$ ). +" However, as discussed earlier, there is a region in the parameter space for the IR solution for which Yy<1."," However, as discussed earlier, there is a region in the parameter space for the IR solution for which $Y_H \le 1$." + In this case only a small fraction of the energy goes into the high energy photons and it is possible (depending on time scales) that most of the electrons cool down rapidly before pair avalanche arises., In this case only a small fraction of the energy goes into the high energy photons and it is possible (depending on time scales) that most of the electrons cool down rapidly before pair avalanche arises. + A natural source of the seed photons is synchrotron emission by the same electrons that produce the IC emission., A natural source of the seed photons is synchrotron emission by the same electrons that produce the IC emission. + Assuming that this source is indeed synchrotron we can proceed and estimate the strength of the magnetic filed and the size of the emitting region., Assuming that this source is indeed synchrotron we can proceed and estimate the strength of the magnetic filed and the size of the emitting region. + We can then check if thesevalues are reasonable within given GRB models., We can then check if thesevalues are reasonable within given GRB models. +" However, we choose"," However, we choose" +respect to the analysis of JSL2).,respect to the analysis of JSL. +" For the SC-tvpe stars. it is a combination of that. the use of a different grid of model atmospheres ancl. in the case of CY Cvg and GP Ori. the adoption here of à T,p;~300 Ix cooler."," For the SC-type stars, it is a combination of that, the use of a different grid of model atmospheres and, in the case of CY Cyg and GP Ori, the adoption here of a $_{eff}\sim 300$ K cooler." + It should be mentioned that the structure of the C-vich atinosphliere models change dramatically with a tiny. variation of the C/O ratio when it is verv close to 1. as it is the case for SC stars.," It should be mentioned that the structure of the C-rich atmosphere models change dramatically with a tiny variation of the C/O ratio when it is very close to 1, as it is the case for SC stars." + Therefore. the abundance analvsis in (hese stars ave affected by systematic differences between the model atmosphere and the real star and/or the treatment of the molecular equilibrium in the computation of the atmosphere models.," Therefore, the abundance analysis in these stars are affected by systematic differences between the model atmosphere and the real star and/or the treatment of the molecular equilibrium in the computation of the atmosphere models." + In consequence. for SC-twpe stars the uncertainty in |F/11] is certainly larger than +0.3 dex.," In consequence, for SC-type stars the uncertainty in [F/H] is certainly larger than $\pm 0.3$ dex." + The main consequence of the new F abundances is that the large [F/Fe] (or [F/O]) radios (up to 1.5 dex) found by JSL in Galactic AGB C-stars are svstematically reduced., The main consequence of the new F abundances is that the large [F/Fe] (or [F/O]) ratios (up to 1.8 dex) found by JSL in Galactic AGB C-stars are systematically reduced. + The largest F enhancements are now close to ~1 dex., The largest F enhancements are now close to $\sim 1$ dex. + These enhancements can be accounted for by current low-mass TP-AGB nucleosynthesis models of solar metallicitv. as we will show below.," These enhancements can be accounted for by current low-mass TP-AGB nucleosynthesis models of solar metallicity, as we will show below." + As noted in $1. during the ascension along the AGB. fresh carbon is mixed within ihe envelope due to TDU episodes.," As noted in 1, during the ascension along the AGB, fresh carbon is mixed within the envelope due to TDU episodes." + Eventually. an O-rich AGB star becomes a C-star when the C/O ratio in (he envelope exceeds 1.," Eventually, an O-rich AGB star becomes a C-star when the C/O ratio in the envelope exceeds 1." + Similarly. F is also expected to increase in the envelope during the AGB phase. thus a fInorine vs. carbon correlation should exist.," Similarly, F is also expected to increase in the envelope during the AGB phase, thus a fluorine vs. carbon correlation should exist." + Figure 2 shows the observed relationship derived in (his study., Figure 2 shows the observed relationship derived in this study. + We have also included in this ligure the O-rich. AGB stars studied by JSL (not analvsed here. open circles).," We have also included in this figure the O-rich AGB stars studied by JSL (not analysed here, open circles)." + Excluding the stars (triangles). a clear increase of the F abundance with the C abundance can be seen.," Excluding the stars (triangles), a clear increase of the F abundance with the C abundance can be seen." + This behaviour is well reproduced by theoretical AGB models., This behaviour is well reproduced by theoretical AGB models. + Lines in Fig., Lines in Fig. + 2 show the predicted F and C content in the envelope (Cristallo et al., 2 show the predicted F and C content in the envelope (Cristallo et al. + 2010) for a 1.5 M. with different metallicities (continuous lines) and 2 M. model (dashed line) with solar metallicitv., 2010) for a 1.5 $_\odot$ with different metallicities (continuous lines) and 2 $_\odot$ model (dashed line) with solar metallicity. + These metallicities match those of the stellar sample (—0.5X [Fe/II| 0.1)., These metallicities match those of the stellar sample $-0.5\leq$ $\leq 0.1$ ). + The theoretical curves start with an envelope composition as determined by the first dredge-up and end at the last, The theoretical curves start with an envelope composition as determined by the first dredge-up and end at the last +compilation of the BCS.,compilation of the BCS. + All 107 have measured redshifts of ) 418: 100 of them fall within the redshift limit of the complete sample at ς<0.3.," All 107 have measured redshifts of $z\le 0.418$ ; 100 of them fall within the redshift limit of the complete sample at $z\le +0.3$." + By design. the X-ray fluxes of the clusters in he extension range from 2.8.10. οσο 7s tod. 10.7 erg 7 s lin the 0.1-2.4 keV band (the latter value being the Hux limit of the original BCS).," By design, the X-ray fluxes of the clusters in the extension range from $2.8\times 10^{-12}$ erg $^{-2}$ $^{-1}$ to $4.4\times +10^{-12}$ erg $^{-2}$ $^{-1}$ in the 0.1–2.4 keV band (the latter value being the flux limit of the original BCS)." + Table | lists the 107 clusters in the BCS extension in analogy o Table 3 of Paper I. i.e.. the contents of this table are One of the clusters listed in Table 1. A1758a. was already listed in Table 3 of Paper I because it made the flux limit of the original sample when combined with its X-ray fainter companion A|758b.," Table 1 lists the 107 clusters in the BCS extension in analogy to Table 3 of Paper I, i.e., the contents of this table are One of the clusters listed in Table 1, A1758a, was already listed in Table 3 of Paper I because it made the flux limit of the original sample when combined with its X-ray fainter companion A1758b." + The latter falls below the flux limit of both the original and the extended BCS., The latter falls below the flux limit of both the original and the extended BCS. + The distribution of the full eBCS sample of 310 clusters in luminosity-redshift space is shown in Fig. 2.., The distribution of the full eBCS sample of 310 clusters in luminosity-redshift space is shown in Fig. \ref{ebcs_lum_z}. + While 211 c0) of these are Abell clusters. the number of non-Abell clusters in the eBCS is substantial: 42 )) systems are Zwicky clusters without Abell identification. and another 57 )) are listed in neither of the two largest optical cluster catalogues.," While 211 ) of these are Abell clusters, the number of non-Abell clusters in the eBCS is substantial: 42 ) systems are Zwicky clusters without Abell identification, and another 57 ) are listed in neither of the two largest optical cluster catalogues." + As expected. the Abell content of the eBCS is thus somewhat lower than. but still similar to. the one found for the BCS where the fractional content in Abell. Zwicky and other clusters was measured to be 70. 11. and 19 per cent.," As expected, the Abell content of the eBCS is thus somewhat lower than, but still similar to, the one found for the BCS where the fractional content in Abell, Zwicky and other clusters was measured to be 70, 11, and 19 per cent." + The redshift distribution of the 310 clusters of the extended BCS shows striking signs of large-scale structure as already noted in the original BCS., The redshift distribution of the 310 clusters of the extended BCS shows striking signs of large-scale structure as already noted in the original BCS. + Figure 3 shows the eBCS redshift histogram compared to the distribution expected for a spatially homogenous sample., Figure \ref{ebcs_n_z} shows the eBCS redshift histogram compared to the distribution expected for a spatially homogenous sample. + The pronounced peaks at 2~0.036 and z~O.077. as well as the depletion between them. are prominent also in the redshift distribution of the original BCS (see Section 8.1 of Paper I.," The pronounced peaks at $z\sim 0.036$ and $z\sim 0.077$, as well as the depletion between them, are prominent also in the redshift distribution of the original BCS (see Section 8.1 of Paper I)." + As shown in Paper IL. the excess of clusters at these redshifts can not be attributed to any single supercluster but is generated by clusters and superclusters distributed widely over the solid angle covered by the BCS.," As shown in Paper I, the excess of clusters at these redshifts can not be attributed to any single supercluster but is generated by clusters and superclusters distributed widely over the solid angle covered by the BCS." + An overview of the distribution of the eBCS on the sky is shown in Fig. 4.., An overview of the distribution of the eBCS on the sky is shown in Fig. \ref{ebcs_skymap}. + The authors are indebted to the ROSAT team at. MPE for providing the RASS photon data this analysis is based upon., The authors are indebted to the ROSAT team at MPE for providing the RASS photon data this analysis is based upon. + The identification of non-cluster sources that contaminated the sample we originally started from was greatly facilitated by the availability of digitized optical images from the POSS and UK Schmidt sky, The identification of non-cluster sources that contaminated the sample we originally started from was greatly facilitated by the availability of digitized optical images from the POSS and UK Schmidt sky +For cach observation of NGC300 X-1. we modeled the ph and combined MOS spectra simultaucously. with each model cousisting of a power law | ciuission line. sufferimg absorption aud iuncludiug a coustaut of uormalization to account for differences in pu and MOS calibration.,"For each observation of NGC300 X-1, we modeled the pn and combined MOS spectra simultaneously, with each model consisting of a power law + emission line, suffering absorption and including a constant of normalization to account for differences in pn and MOS calibration." + The resulting best fits are shown in Table 3.., The resulting best fits are shown in Table \ref{specfits}. + We found the spectra from Obs., We found the spectra from Obs. + 2.1 to be consistent with he model reported by ?.., 2–4 to be consistent with the model reported by \citet{carp07a}. + However. the best fit to the Obs.," However, the best fit to the Obs." + 1 spectra vielded ναοί = 100/60. an unacceptable fit.," 1 spectra yielded $\chi^2$ /dof = 100/60, an unacceptable fit." + We show this best fit iu Fie. L.," We show this best fit in Fig. \ref{ob1300spec}," + alone with \? residuals., along with $\chi^2$ residuals. + We see that the model svstematically underestimates the emission at higher energies. suggestive of anu additional emiüssiou component.," We see that the model systematically underestimates the emission at higher energies, suggestive of an additional emission component." + IHeuce. we fitted all NGC300. X-1 spectra with absorbed: blackbody power law models aud also absorbed dise blackbody | power law models. commonly observed in NS and BIL NBs respectively.," Hence, we fitted all NGC300 X-1 spectra with absorbed blackbody + power law models and also absorbed disc blackbody + power law models, commonly observed in NS and BH XBs respectively." + The best fits fx these models are also provided in Table 3::3: 0.3.10 keV Iuimuinosities are provided. assuniug a distance of1.58 Mpc (2)..," The best fits for these models are also provided in Table \ref{specfits}; 0.3–10 keV luminosities are provided, assuming a distance of 1.88 Mpc \citep{gps05}." + We see that both two-component models provided eood fits to the Obs., We see that both two-component models provided good fits to the Obs. + 1 aud. Obs., 1 and Obs. + 3 spectra., 3 spectra. + The Obs., The Obs. + 3 spectrum is well described by all three models: lOWOCVOT. the resulting bhunuinosities differ bv factor of two (~ha): hence theuncertain.," 3 spectrum is well described by all three models; however, the resulting luminosities differ by factor of two $\sim$ $\sigma$ ); hence the." + Neither the Obs., Neither the Obs. + 2 nor Obs., 2 nor Obs. + | spectra were successfully described by any thermal | power law model: hence NCC3000 X-L exhibits at least two spectra states: one exhibited in Obs., 4 spectra were successfully described by any thermal + power law model; hence NGC3000 X-1 exhibits at least two spectral states: one exhibited in Obs. + 1 aud the other exhibited iu Obs., 1 and the other exhibited in Obs. + 2 and 1., 2 and 4. + Finally. we fitted the spectra withstar.," Finally, we fitted the spectra with." + Our literature review shows that 1ο wind-accreting DAINBs with PDS similar to those observed from NCGC300. N-1 all have spectra that are well lescribed by a power law with Fo<1.8., Our literature review shows that the wind-accreting HMXBs with PDS similar to those observed from NGC300 X-1 all have spectra that are well described by a power law with $\Gamma$ $\le$ 1.8. + Hence. we modeled 16 spectra from cach observation with power law models. restricting-. P to be «1.5.," Hence, we modeled the spectra from each observation with power law models, restricting $\Gamma$ to be $<$ 1.8." +E* We- foundJ the best fitd X2g /doft to f 25 qm each case., We found the best fit $\chi^2$ /dof to be $\ga$ 5 in each case. + Hence δις200 X-1 is unlikely to be vowered by Boudi-Hovle accretion outo a neutron star., Hence NGC300 X-1 is unlikely to be powered by Bondi-Hoyle accretion onto a neutron star. +" We first attempted to sumnultaucouslv fit the pu aud combined MOS spectra of ICLO N-1 with various ciission models, but failed to obtain a good fit."," We first attempted to simultaneously fit the pn and combined MOS spectra of IC10 X-1 with various emission models, but failed to obtain a good fit." + This is likely to be due to the fact that ICLO X-1 varies in colour aud iuteusitv ching the observation. mixing the spectral states so that a plivsical ft is impossible.," This is likely to be due to the fact that IC10 X-1 varies in colour and intensity during the observation, mixing the spectral states so that a physical fit is impossible." +time constraints.,time constraints. + For all sources. the spectrometer was centered on the velocity of peak cuuission of the water maser.," For all sources, the spectrometer was centered on the velocity of peak emission of the water maser." + However. the wide velocity coverage precludes the possibility of iiethanol masers beue missed due to their volocities being significantly offset from that of the water lasers.," However, the wide velocity coverage precludes the possibility of methanol masers being missed due to their velocities being significantly offset from that of the water masers." + Our observations resulted in the discovery of 10 new nethanol masers. the properties of which are listed iu Table 2.," Our observations resulted in the discovery of 10 new methanol masers, the properties of which are listed in Table 2." + Since we did not attempt to refine hne position of the methanol masers using a evid of observations. he positious quoted in Table 2 could have an error as Πο] as ~LW.," Since we did not attempt to refine the position of the methanol masers using a grid of observations, the positions quoted in Table 2 could have an error as high as $\sim 1'$." + The maser phuuinosities quote in Table 2 are calculated from the iutegrated flux deusity assuming isotropic enissiou., The maser luminosities quoted in Table 2 are calculated from the integrated flux density assuming isotropic emission. + Details of water maser sources that hac rou-detectious of 1iethanol masers are indicated in Table 3. which is available on-line.," Details of water maser sources that had non-detections of methanol masers are indicated in Table 3, which is available on-line." + The peak flux deusities of he inethanol mascrs detected in our survey range from 5 to 0 Jy., The peak flux densities of the methanol masers detected in our survey range from 0.5 to 4.0 Jy. + Five sources are located bevoud the solar circle. which is a significant additiou to the παπα of such sources in the northern sky.," Five sources are located beyond the solar circle, which is a significant addition to the number of such sources in the northern sky." + It is interesting that the sineniatic distance o the source 051257|3919 puts it a a distance of ~ Ll kpc from the Sun aud. 20 spe frou he Calactic center., It is interesting that the kinematic distance to the source 05137+3919 puts it at a distance of $\sim$ 14 kpc from the Sun and 20 kpc from the Galactic center. + This is oue of the farthes nethanol llasors (in terms of distance from the Calactic ceuter) in the outer Calaxw. although the uucertaiuties in the rotation curve at these galactroceutriic radii aud peculiar notions such as tha observed in W3OII (Xu et al.," This is one of the farthest methanol masers (in terms of distance from the Galactic center) in the outer Galaxy, although the uncertainties in the rotation curve at these galactrocentric radii, and peculiar motions such as that observed in W3OH (Xu et al." + 2006) ranslae to significant uncertainties in the kincmatic distance., 2006) translate to significant uncertainties in the kinematic distance. + Measiiug parallax distances to sources like 05137|3919. will be useful for measuring the rotation speed of the Calaxy at large ealactoceutric rac., Measuring parallax distances to sources like 05137+3919 will be useful for measuring the rotation speed of the Galaxy at large galactocentric radii. + The spectra of the CIT4OIT inasers detected in our survey are shown iu Figure 1., The spectra of the $_{3}$ OH masers detected in our survey are shown in Figure 1. + The spectra have a velocity resolution of 0.11 kan !., The spectra have a velocity resolution of 0.11 km $^{-1}$. + Here we present notes on individual sources., Here we present notes on individual sources. +-2-3919.. TThere are two features that are separated by about [5 lan 1., There are two features that are separated by about 4.5 km $^{-1}$. + The stronger feature is at an LSR velocity of 3.9 1n +. while the weaker feature has a flux density of only 0.35 Jv.," The stronger feature is at an LSR velocity of –3.9 km $^{-1}$, while the weaker feature has a flux density of only 0.35 Jy." + This region is associated with a 3.6 ci continuunui source (Molinari et al., This region is associated with a 3.6 cm continuum source (Molinari et al. + 2002)., 2002). +06446-4-0029.. There are at least five features over a velocity range of over 7 kn +., There are at least five features over a velocity range of over 7 km $^{-1}$. + The feature at 18.6 lau | is the strongest oue., The feature at 48.6 km $^{-1}$ is the strongest one. + There is wear infrared cutission in this region as seen in 2ATASS. aud a non-detection of SiO iasers (Harju et al.," There is near infrared emission in this region as seen in 2MASS, and a non-detection of SiO masers (Harju et al." + 1005).18319-0802., 1998). +. TThere are several features spauniug a velocity range over 15 kins +., There are several features spanning a velocity range over 15 km $^{-1}$. + The weakest feature is ouly about 0.2 Jv., The weakest feature is only about 0.2 Jy. + Au ultracompact (UC) HIE region. separated by about LO”. could be associated with this region (Becker ot al.," An ultracompact (UC) HII region, separated by about $''$, could be associated with this region (Becker et al." + 1991)., 1994). +18355-0650.. T'There are at leas five features spauline a velocity range of about 6 kin 1, There are at least five features spanning a velocity range of about 6 km $^{-1}$. +" Au UC UII region. separated by 17"", is associated with this resion (Becker et al."," An UC HII region, separated by $''$, is associated with this region (Becker et al." + 1991)., 1994). + Features span from 93.1 to 105.3 lan |! with multiple features being blended toecther., Features span from 93.4 to 105.3 km $^{-1}$ with multiple features being blended together. + There are at least two compact TTD regions associated with this region within 2’ (Wood Clawcluwell 19892: Decker et al., There are at least two compact HII regions associated with this region within $'$ (Wood Churchwell 1989a; Becker et al. + 199)., 1994). +" One of them. separated wv about SO"", could be associated with this region."," One of them, separated by about $''$, could be associated with this region." +18403-0440.. TThis source primarily shows a single feature at 20.1 kin, This source primarily shows a single feature at +20.1 km $^{-1}$. + No «servations have been reported on this region. except near mfrared emission from the 2MASS.," No observations have been reported on this region, except near infrared emission from the 2MASS." +18479-0005.. Walsh ct al. (, Walsh et al. ( +1997) did not detect mascr Cluission within them 1 Jw limit (80).,1997) did not detect maser emission within their 1 Jy limit $\sigma$ ). + There are several features crowded within a velocity range of ouly 2 kimi 1, There are several features crowded within a velocity range of only 2 km $^{-1}$. + A UC ΠΠ region is associated with this source (IXurtz et al., A UC HII region is associated with this source (Kurtz et al. + 1991)., 1994). +--4001.. TThis source displavs a single feature at -6.9 kan +. which matches he II?CO! (1-3) peak of Tasceawa Mitchell. (1995).," This source displays a single feature at -6.9 km $^{-1}$, which matches the $^{13}$ $^{+}$ (4-3) peak of Hasegawa Mitchell (1995)." + A bipolu outflow aud a continmun source were also detected im this region (asceawa Alitchell 1995: Trinidad et al., A bipolar outflow and a continuum source were also detected in this region (Hasegawa Mitchell 1995; Trinidad et al. + 2003)., 2003). + Alid-iufraxrec inages show that this source is surrounded by au optically thick dusty cuvelope (Mareugo et al., Mid-infrared images show that this source is surrounded by an optically thick dusty envelope (Marengo et al. + 2000)., 2000). +21306+5540.. There are clearly five features in this source., There are clearly five features in this source. + Except for uear infrared euissiou from the 24ASS. uo other observations have reported iu this region.," Except for near infrared emission from the 2MASS, no other observations have reported in this region." +"(8140). TThis source shows a double peaked structure aud is the weakest source detected in tliis survey,", This source shows a double peaked structure and is the weakest source detected in this survey. + At a clistance of 910 pe (Craunpton Fisher 1971). the maser luminosity is verv low (1.1 ?L.).," At a distance of 910 pc (Crampton Fisher 1974), the maser luminosity is very low (1.1 $^{-9} L_{\odot}$ )." + However. its infrared buudinositv mdicates that it is still associated with a massive star forming region (Table 2).," However, its infrared luminosity indicates that it is still associated with a massive star forming region (Table 2)." + A faint coutiuuuni source (liurtz et al., A faint continuum source (Kurtz et al. + 199) aud a CO outflow (Minchiu et al., 1994) and a CO outflow (Minchin et al. + 1993) are associated with the region., 1993) are associated with the region. + We detected 10 nethanol masers frou argetiug 89 sources which results dji a detection rate of ~., We detected 10 methanol masers from targeting 89 sources which results in a detection rate of $\sim$. + However. to compare the statistics of Πο iux methanol Lasers. we have to coisider the cutive orieinal sample of 175 sources that saisfied our selection criteria (0> 107).," However, to compare the statistics of $_2$ O and methanol masers, we have to consider the entire original sample of 178 sources that satisfied our selection criteria $\delta \geq -10^{\circ}$ )." + lr sources have previous detections of methanol mascrs. while 12 sources were not observed due to coustrauts of observiug time.," 47 sources have previous detections of methanol masers, while 42 sources were not observed due to constraints of observing time." + Heuce. the overall detection rate of methanol masers in à water maser sample is at least ~©.," Hence, the overall detection rate of methanol masers in a water maser sample is at least $\sim$." + Our sample includes 17 sources that are associated with low-nass YSOs aud 10 of them were observed iu this survey., Our sample includes 17 sources that are associated with low-mass YSOs and 10 of them were observed in this survey. + A il detection rate from these sources adds to the results of Minier et al. (, A nil detection rate from these sources adds to the results of Minier et al. ( +2003). sueeesting that 6.7 CGIIz methanol masers are only associated with massive star foriung regions.,"2003), suggesting that 6.7 GHz methanol masers are only associated with massive star forming regions." + Therefore. excliding the low-iass YSOs. the detection rate is at least ~ 3554...," Therefore, excluding the low-mass YSOs, the detection rate is at least $\sim$ ." +Ixepleriau angular frequency is of the order of 300 sec.,Keplerian angular frequency is of the order of 300 sec. + This vields the characteristic decay time of order of a week to a mouth., This yields the characteristic decay time of order of a week to a month. + A time scale of this order secs to be right for our model as the repetition time of ruini-outbursts is about a week.," A time scale of this order seems to be right for our model, as the repetition time of mini-outbursts is about a week." + Ou the other mand. once the disk becomes hot aud ionized. the “Balbus-Tawlev instability becomes effective and starts to buik up the maeuetic fields.," On the other hand, once the disk becomes hot and ionized, the 'Balbus-Hawley instability' becomes effective and starts to build up the magnetic fields." + Their erowth iue is thought to be very short. estimated to be of the order of the inverse of the Keplerian rotation frequency. Terowth~Lf/Q.," Their growth time is thought to be very short, estimated to be of the order of the inverse of the Keplerian rotation frequency, $\tau_{\rm growth} \sim +{1/\Omega}$." + This time-scale is sufficicutly short that he decaved magnetic fields can be brought back to the original level aud that thus the viscosity is refreshed to he original high value every time when a mini-outhurst occurs., This time-scale is sufficiently short that the decayed magnetic fields can be brought back to the original level and that thus the viscosity is refreshed to the original high value every time when a mini-outburst occurs. + This large disparity of growth rate and decay rate of naguctic fields is a main inerccdicnt of our model., This large disparity of growth rate and decay rate of magnetic fields is a main ingredient of our model. + We now oxeseut the sinulatious based ou these ideas., We now present the simulations based on these ideas. + We simulate the light curve of the six qaiui-outhbiursts and their sudden cessation in EC Cie by means of a siupli&ed model used for the simulation of the leht curves of SU UMa stars (Osaki 1989) and WZ See (Osalsi 1995a)., We simulate the light curve of the six mini-outbursts and their sudden cessation in EG Cnc by means of a simplified model used for the simulation of the light curves of SU UMa stars (Osaki 1989) and WZ Sge (Osaki 1995a). + Our uodoel is based ou a formulation by Audersou (19088) to treat time evolution of the disk radius in a dwarf nova outburst evele., Our model is based on a formulation by Anderson (1988) to treat time evolution of the disk radius in a dwarf nova outburst cycle. + Tt was used exteusivolv by one of the authors to simulate various licht curves of SU UNa-type dwarf novae (Osaki 1989 for SU UMa stars in eeneral. Osaki 1995a for WZ Sec. Osaki 1995b for ER UMa. Osaki 1995c for RZ LAG).," It was used extensively by one of the authors to simulate various light curves of SU UMa-type dwarf novae (Osaki 1989 for SU UMa stars in general, Osaki 1995a for WZ Sge, Osaki 1995b for ER UMa, Osaki 1995c for RZ LMi)." + The model disk cousists of an inviscid disk component aud a torus conrponeut., The model disk consists of an inviscid disk component and a torus component. + Since the disk component iu our simplifed model is taken as inviscid. ouly au outside-in type outburst cau be treated.," Since the disk component in our simplified model is taken as inviscid, only an outside-in type outburst can be treated." + The outsice-in outburst is triggered when the OrUus lnass exceeds its critical value., The outside-in outburst is triggered when the torus mass exceeds its critical value. + We modify our previous calculations for the WZ See (Osaki 1995a) now takiug into account a viscosity decay. as described in the previous section., We modify our previous calculations for the WZ Sge (Osaki 1995a) now taking into account a viscosity decay as described in the previous section. +" The binary parameters used for EC Cuc in our siuulations are those used for WZ See iu the previous work (Osaki 1995a): (1) the mass of the primary white dwirf. M4=LOAM.: (2) the radius of the white dwarf. πι=0.6< 10am: (3) the mass of the secondary star, AL,=(LAL. (1) the binary separation. A=£L4G:1027013: (5) the cirenlarization radius or Lubow-Shu radius. (0.23.4.=10101."," The binary parameters used for EG Cnc in our simulations are those used for WZ Sge in the previous work (Osaki 1995a): (1) the mass of the primary white dwarf, $M_1=1.0 M_{\odot}$; (2) the radius of the white dwarf, $R_1=0.6\times 10^9$ cm; (3) the mass of the secondary star, $M_2=0.1 M_{\odot}$; (4) the binary separation, $A=4.46\cdot 10^{10}$ cm; (5) the circularization radius or Lubow-Shu radius, $r_{\rm LS}=0.23A=10^{10}$ cm." + We take a mass trauster rate 1.5.lores |., We take a mass transfer rate $\dot M=1.5 \cdot 10^{15}$ $^{-1}$. + Tn order to explain the long quiesceuce of EC Cue. 19 vers. we assume that the quiesceut disk viscosity is extremely low. ρα~0.001.," In order to explain the long quiescence of EG Cnc, 19 years, we assume that the quiescent disk viscosity is extremely low, $\alpha_{\rm cold} \sim 0.001$ ." + In our simplified model. the viscosity appears through the quantity 7. where +; is the torus radius. nes the radius of the disk's outer edge and Ar the radial width of the torus.," In our simplified model, the viscosity appears through the quantity $\beta$, where $r_t$ is the torus radius, i.e., the radius of the disk's outer edge and $\Delta r$ the radial width of the torus." + μις Is the critical masini surface deusitv aud X44 the surface density of the cold disk just after the cud of the outburst. estimated to be Mego&2«Minin (for mere details see Osalà 1989).," $\Sigma_{\rm max}$ is the critical maximum surface density and $\Sigma_{\rm cool}$ the surface density of the cold disk just after the end of the outburst, estimated to be $\Sigma_{\rm +cool} \simeq 2\times \Sigma_{\rm min}$ (for more details see Osaki 1989)." + Therefore. a low viscosity value. a. corresponds to a large X45 and therefore to a large Jj.," Therefore, a low viscosity value, $\alpha$, corresponds to a large $\Sigma_{\rm max}$ and therefore to a large $\beta$." + Tt mw be noted here that we used; =0.6 to simmlate VW Ibi (Osalki 1989) while 3=10 was used to sinulate re long quiescence of WZ Sec (Osaki 1995a)., It may be noted here that we used $\beta=0.6$ to simulate VW Hyi (Osaki 1989) while $\beta=40$ was used to simulate the long quiescence of WZ Sge (Osaki 1995a). + Similarly to explain the long quiescence of EC Cie we adopted +=20 or the pre-outburst cold disk., Similarly to explain the long quiescence of EG Cnc we adopted $\beta=20$ for the pre-outburst cold disk. + The light curve of the main outburst of EG Cuc is calculated in a similar way as that or WZ See (Osaki 19953)., The light curve of the main outburst of EG Cnc is calculated in a similar way as that for WZ Sge (Osaki 1995a). + For the post-outhurst rebrightenine of EC Cuc. we iow Incorporate the viscosity decay clescribed in the xevious section into our simplified model.," For the post-outburst rebrightening of EG Cnc, we now incorporate the viscosity decay described in the previous section into our simplified model." + Ta order to accommodate the exponential decay of viscosity. we take he quantity 9 not constant in time. but exponcutially increasing in the followine wav where oy Is an initial value of 0. and 7 ix a characteristic time for the viscosity decay. of the order of aweek to aiouth. aud £ is the time counted from the eud of either the main outburst or cach iuiui-outburst.," In order to accommodate the exponential decay of viscosity, we take the quantity $\beta$ not constant in time, but exponentially increasing in the following way where $\beta_0$ is an initial value of $\beta$, and $\tau$ is a characteristic time for the viscosity decay, of the order of a week to a month, and $t$ is the time counted from the end of either the main outburst or each mini-outburst." + That Is. Viscosity is assinued to be refresliec to a high value each time a miui-outburst occurs.," That is, viscosity is assumed to be refreshed to a high value each time a mini-outburst occurs." + In this paper. the quantities à and 7 are two adjustable parameters to reproduce the sequence of rapid 1uiui-outbursts with a mean recurrence time of seven davs.," In this paper, the quantities $\beta_0$ and $\tau$ are two adjustable parameters to reproduce the sequence of rapid mini-outbursts with a mean recurrence time of seven days." + This indicates a fairlyhigh viscosity and requires jy to be small., This indicates a fairlyhigh viscosity and requires $\beta_0$ to be small. + We have varied the parameter values for these two quantities and found that ο=0.08 aud 7=25 davs, We have varied the parameter values for these two quantities and found that $\beta_0=0.08$ and $\tau=25$ days +corresponding to a wavelength A=2x/kz10 m. For &czAy. the erowth rate from eq. (48)),"corresponding to a wavelength $\lambda=2\pi/k\simeq 10$ m. For $k>>k_c$, the growth rate from eq. \ref{highk}) )" + is The hvdrocdynamic treatment is restricted to ο<\Omega$." + I vortex creep is in the stronglv-dampoed regime 37a. contrary to the estimates here. there is still à broad. window for instability.," If vortex creep is in the strongly-damped regime $\beta>>\alpha$, contrary to the estimates here, there is still a broad window for instability." + Requiring Ay«Fs gives and the star will be unstable at some wavenumber that is consistent. with the hvedrodynamic approximation er<<ο)., Requiring $k_c1 is an integer.," Therefore, for large $z$, $p(z)\sim\exp(-\alpha z^{2m})$ , where $\alpha$ is a positive constant and $m\ge 1$ is an integer." + With these provisos. db turns out that for all plausible frequency [unctions. ∖∖⊽≼↲∪∐≱≼↲↕⋅↥∖∖⊽∪⊳∖⇁↕∐↓↕↽≻↥≼↲≼↲⇀↸≀↧↴∐↓↕↽≻↥≼↲⋝∖⊽⋃⋟∪↕⋅∖∖↽∐↕≺∢∐∣∣∣∶⊔∶ These are merely examples.," With these provisos, it turns out that for all plausible frequency functions, We offer two simple examples (for which$m=1$ ): These are merely examples." + The exact value for Q is not important for this study., The exact value for $Q$ is not important for this study. + Eq., Eq. + is entirely. aclequate to restate the condition for equilibrium. Eq. 22..," \ref{qap} is entirely adequate to restate the condition for equilibrium, Eq. \ref{qr}," + as The 15 components of the HF disk mass model were itemized byFlynnεἰal. 2006.., as The 15 components of the HF disk mass model were itemized by\citealt{Fetal:06}. . + From these.we caleulate," From these,we calculate" +The ciseretization scheme for the non-relativistic uid equations is summarised in Figure Al.,The discretization scheme for the non-relativistic fluid equations is summarised in Figure \ref{fig:scheme}. + Fluxes are caleulatecl on the, Fluxes are calculated on the +fields are higher by ~20 kiu/s for all pp<112.,fields are higher by $\sim 20$ km/s for all $\mu<14.2$. + Also note the difference between the star counts with gp between the three models., Also note the difference between the star counts with $\mu$ between the three models. + Velocities iu the boxy bulge are outside the scope of this paper: they have already been studied with a similar N-body simulation by Shenetal.(2010). and compared with results from the BRAVA survey., Velocities in the boxy bulge are outside the scope of this paper; they have already been studied with a similar N-body simulation by \citet{Shen+10} and compared with results from the BRAVA survey. + We have analyzed star counts in the immer Galaxy using au N-body model which arose from secular evolution of a disk ealaxy. a natural mechanisin for galaxies like ours.," We have analyzed star counts in the inner Galaxy using an N-body model which arose from secular evolution of a disk galaxy, a natural mechanism for galaxies like ours." + The bar and buckling iustabilities in the stellar disk lead to a boxy bulee which exteuds to a longer in-plane bar., The bar and buckling instabilities in the stellar disk lead to a boxy bulge which extends to a longer in-plane bar. + The bar couples with the spiral avis in the disk. ceiving rise alternately to leading. straight or trailing bar euds.," The bar couples with the spiral arms in the disk, giving rise alternately to leading, straight or trailing bar ends." + As seen from within the disk at 8 kpe from the center. he mmaxima of the line-ofsight distance distributions in the Calactic plane occur at distances somewhat urther than the maxima of the line-ofsight density distributions. due to the voluue effect in the star counts.," As seen from within the disk at 8 kpc from the center, the maxima of the line-of-sight distance distributions in the Galactic plane occur at distances somewhat further than the maxima of the line-of-sight density distributions, due to the volume effect in the star counts." + Asstuning a plausible oricutation (a=25°). this explains xut of the observational signature which was previously used to infer the existence of a secondbar.," Assuming a plausible orientation $\alpha=25^\circ$ ), this explains part of the observational signature which was previously used to infer the existence of a second." + If in addition we choose a model suapshot where the bar has cading ends. most of thebar signature iu the star count data can be reproduced.," If in addition we choose a model snapshot where the bar has leading ends, most of the signature in the star count data can be reproduced." + While not made specially o 1match the MW. this model thus illustrates that the raditional Galactic bar (the boxy bulee) aud the more recently inferredbar can plausibly be explained by a singlebar structure.," While not made specially to match the MW, this model thus illustrates that the traditional Galactic bar (the boxy bulge) and the more recently inferred can plausibly be explained by a single structure." + To test this further we have determined the dependence of the mean racial velocity and velocity dispersiou on distance inodulus in Calactic plane fields σα the inferred end of the planar bar., To test this further we have determined the dependence of the mean radial velocity and velocity dispersion on distance modulus in Galactic plane fields near the inferred end of the planar bar. + These illustrate the differcuces between the barred model and an axisviunietric rotating disk. which cau be compared witli upconiug radial velocity survey data for the immer MW.," These illustrate the differences between the barred model and an axisymmetric rotating disk, which can be compared with upcoming radial velocity survey data for the inner MW." + Future work should also address in nore detail the spiral arm - bar interaction which gives rise to the curved euds of the bar. and aiu at coustructing a more detailed dynamical model allowing us to uuderstaud better the structure aud evolution of the immer Milky. Wav.," Future work should also address in more detail the spiral arm - bar interaction which gives rise to the curved ends of the bar, and aim at constructing a more detailed dynamical model allowing us to understand better the structure and evolution of the inner Milky Way." + We thauk wen Freeman for helpful discussions about this work., We thank Ken Freeman for helpful discussions about this work. + IMV thauks Peter Tauunerslev for discussions ou the MW., IMV thanks Peter Hammersley for discussions on the MW. +A striking feature of the velocitv-position maps with nupact parameters greater than about 20 arcsec is the appearance of an clliptical shape for the outermost cluission (Fig. L.,A striking feature of the velocity-position maps with impact parameters greater than about 20 arcsec is the appearance of an elliptical shape for the outermost emission (Fig. \ref{fig4}. + Eniission is very weak. possibly absent. bevond about 1 minute of are from the star.," Emission is very weak, possibly absent, beyond about 1 minute of arc from the star." + Ellipticity of isophotes iu the velocity-position map corresponds to a spherically Μποτς expaudiug shell., Ellipticity of isophotes in the velocity-position map corresponds to a spherically symmetric expanding shell. + Geometry of the slit position (at 33 aresec from the star) and the seni-major axis of the ellipse show that the shell radius is 50 + 2 aresec., Geometry of the slit position (at 33 arcsec from the star) and the semi-major axis of the ellipse show that the shell radius is 50 $\pm$ 2 arcsec. + The ανα velocity width. the cllipse’s secnmundünor axis corresponds to an expansion (radial) velocity of 18 c 2lau |.," The maximum velocity width, the ellipse's semiminor axis, corresponds to an expansion (radial) velocity of 18 $\pm$ 2km $^{-1}$." + Note that the shell is clearly seen in Fig., Note that the shell is clearly seen in Fig. + 5 as the plateau followed by au intensity drop at about 50 arcsec., \ref{fig6} as the plateau followed by an intensity drop at about 50 arcsec. + A that distance from the star. the mean deusity is rapidly decreasing with distance. due to the interstellar UV. field (see. Fig.," At that distance from the star, the mean density is rapidly decreasing with distance, due to the interstellar UV field (see Fig." + 5 aud Sec. ??)).," \ref{fig6} + and Sec. \ref{radial}) )," + aud the shell represeuts a considerable deusity enlianceioent., and the shell represents a considerable density enhancement. + The hickness of the shell is 1 to 2 aresec in position ou the sky and is wuresolved in velocity. ic. less than 2.6 kin +)," The thickness of the shell is 1 to 2 arcsec in position on the sky and is unresolved in velocity, i.e. less than 2.6 km $^{-1}$." + Non-uniform euission over the spherical shell could result from one or more of several factors: shell ejection occurred prefercutially over one hemisphere of the star: the excitation of the shell is non-uniform owing to a bright spot on the stellar disk: the ionisation of potassium atoms is ereater over one hemisphere due to a chromosphleric enhancement: the shell is plowing iuto the nregular local interstellar medium and deceleration euliances the local density non-uniforunily., Non-uniform emission over the spherical shell could result from one or more of several factors: shell ejection occurred preferentially over one hemisphere of the star; the excitation of the shell is non-uniform owing to a bright spot on the stellar disk; the ionisation of potassium atoms is greater over one hemisphere due to a chromospheric enhancement; the shell is plowing into the irregular local interstellar medium and deceleration enhances the local density non-uniformly. + Circuustellar absorption imposed upon the plotospheric spectrum is composed of tivo components in atomic (ce. T699 aas shown Ez Bernat Lambert ΠΟΤΟ and Goldberg ct al. [," Circumstellar absorption imposed upon the photospheric spectrum is composed of two components in atomic (e.g., 7699 as shown by Bernat Lambert [1975] and Goldberg et al. [" +1975 xb the CO lines by Bernat et al. (,1975]) and the CO lines by Bernat et al. ( +1979) who labelled the components S1 and $2 where the former has a higher excitation temperature (200Í vs TOW). substantially larger CO and column densities and lower expansion velocities.,"1979) who labelled the components S1 and S2 where the former has a higher excitation temperature (200K vs 70K), substantially larger CO and column densities and lower expansion velocities." + The expansion velocity deduced. frou the $2 absorption lines (20.2 kin |! from Goldberg et ab.," The expansion velocity deduced from the S2 absorption lines (20.2 km $^{-1}$ from Goldberg et al.," + aud 13 kui ! from Bernat ct al), and 13 km $^{-1}$ from Bernat et al.) + are consistent with our measurement., are consistent with our measurement. + Note that the absorption Lue measurements are made relative to photospheric lines but the photospheric velocity is variable., Note that the absorption line measurements are made relative to photospheric lines but the photospheric velocity is variable. + Our measurement is a true expansion velocity. relative to the ceuter-ofimass of the star. indepeucert of the stellar velocity. and of plotospheric velocity variations.," Our measurement is a true expansion velocity, relative to the center-of-mass of the star, independent of the stellar velocity, and of photospheric velocity variations." + Another notable difference between he S1 aud $2 conmponents is the very narrow line width of the latter: a Doppler width of 1 kins ! for S2 but Eus | for SI roni visual lines (Bernat 1977)., Another notable difference between the S1 and S2 components is the very narrow line width of the latter: a Doppler width of 1 km $^{-1}$ for S2 but 4 km $^{-1}$ for S1 from visual lines (Bernat 1977). + The Doppler width for S2 corresponds to a full width at half iuteusitv of 1.7 kii loan estimate consistent with the upper Iit obtained roni our observation of the shell emissiou.," The Doppler width for S2 corresponds to a full width at half intensity of 1.7 km $^{-1}$, an estimate consistent with the upper limit obtained from our observation of the shell emission." + Bernat et al., Bernat et al. + roni a conrparison of the collisional aud radiative rates or CO excitation inferred that the former dominate aud. rence. them excitation teniperature was approxinatelv equivalent to the eas kinetic temperature.," from a comparison of the collisional and radiative rates for CO excitation inferred that the former dominate and, hence, their excitation temperature was approximately equivalent to the gas kinetic temperature." + Boference to a dust shell model (Tsuji 1979) showed a gas eni])oratire of 70 TI& was achieved at about 55 arcsec from the star. a prediction quite consistent with our observed shell radius.," Reference to a dust shell model (Tsuji 1979) showed a gas temperature of 70 K was achieved at about 55 arcsec from the star, a prediction quite consistent with our observed shell radius." + Judged by velocity of expansion. line width. aud shell radius as measured or estimated from the CO lines aud our yvelocity-position maps. we identity the circumstellar gas contributing the $2 absorptio- σαines with the outermost lin shell seen iu Fig. L.," Judged by velocity of expansion, line width, and shell radius as measured or estimated from the CO lines and our velocity-position maps, we identify the circumstellar gas contributing the S2 absorption lines with the outermost thin shell seen in Fig. \ref{fig4}." + This ideuti&cation allows interesting further deductious to be made., This identification allows interesting further deductions to be made. + Ii particular. he volume density of CO molecules aud neutral Is atoms nav be estimated from their column deusity anc he shell thickness.," In particular, the volume density of CO molecules and neutral K atoms may be estimated from their column density and the shell thickness." + The equivalent width of the S2 line estimated from Goldberg et al. (, The equivalent width of the S2 line estimated from Goldberg et al. ( +1975. their Figure ]a) gives a colhuun deusitv DomLS «1013 on the assumption that the line is unsaturated.,"1975, their Figure 1a) gives a column density ) $\simeq +1.3 \times 10^{11}$ $^{-2}$ on the assumption that the line is unsaturated." + After correction for sccing. the shell’s thickness may be about 1 second of arc.," After correction for seeing, the shell's thickness may be about 1 second of arc." + At d = 10 pe. a shell thickness of 1 second of are corresponds to 2: «1057 cm.," At $d$ = 140 pc, a shell thickness of 1 second of arc corresponds to 2 $\times 10^{15}$ cm." + Then. the density in the shellis about 6 «10. 7.," Then, the density in the shell is about 6 $\times 10^{-5}$ $^{-3}$." + In contrast. Rodgers Classgold's (1991) model predicted a density of μεν 1)2«109 P at the shells distance.," In contrast, Rodgers Glassgold's (1991) model predicted a density of $n$ $\simeq 2 + \times 10^{-6}$ $^{-3}$ at the shell's distance." + This difference is not surprising as the shell by inspection of, This difference is not surprising as the shell by inspection of +does indeed exist for a reasonable set of parameters.,does indeed exist for a reasonable set of parameters. +" To derive an order of magnitude estimate of the possible scatter, for each of the parameters we estimated the range into which, leaving the other parameters unchanged, one still obtains E>“=100 MeV, Emaxz1 GeV, and a flux level at 1 GeV compatible with the LAT observations (green vertical line in Fig."," To derive an order of magnitude estimate of the possible scatter, for each of the parameters we estimated the range into which, leaving the other parameters unchanged, one still obtains $E^{SC}_{p} \gtrsim 100$ MeV, $E_{max} \gtrsim 1$ GeV, and a flux level at 1 GeV compatible with the LAT observations (green vertical line in Fig." + 2)., 2). +" In this way, we obtain pz2.7, 0.8x10?€Ls.2x102, 0.42€&0.5 (where we set the upper limit to ensure that not more than half of the internal energy goes into accelerating the electrons), 2X10?0$, more emission is predicted at negative velocities at low latitudes and $l=180^\circ$." + Also less emission is predicte at negative velocities and |/|=-10 and [οz30, Also less emission is predicted at negative velocities and $|l|\simeq\pm10^\circ$ and $|b|\simeq30^\circ$. + oth these changes arise from an increase in the extent of global inflow of theILL. and they improve the fit to the data.," Both these changes arise from an increase in the extent of global inflow of the, and they improve the fit to the data." + Moreover. with az0 the optimum value of. fi; is reduced from unity to 0.3. so clouds become neutral abou30 percent of the way to their highest point above the plane.," Moreover, with $\alpha\ne0$ the optimum value of $f_{\rm ion}$ is reduced from unity to $0.3$ , so clouds become neutral about$30$ percent of the way to their highest point above the plane." + Our best-litting model has a=6.8Gyr which implies," Our best-fitting model has $\alpha=6.3\Gyr^{-1}$ , which implies" +amplitude must go through a minimum at some epoch.,amplitude must go through a minimum at some epoch. + For low mass galaxies. this minimum could be at z>0.," For low mass galaxies, this minimum could be at $z \gg 0$." + The assumption outlined above applies in a much better way to quasars as these cannot form in very low mass halos anc the redshift at which the amplitude of quasar correlation reaches the minima may be low enough to be observable., The assumption outlined above applies in a much better way to quasars as these cannot form in very low mass halos and the redshift at which the amplitude of quasar correlation reaches the minima may be low enough to be observable. + LE 10 minimum. halo mass associated: with quasars is greater van 1024A7. then the observed. quasar correlation function garoulcl have a higher amplitude at. z2 than at z«I. ο., If the minimum halo mass associated with quasars is greater than $10^{11} M_\odot$ then the observed quasar correlation function should have a higher amplitude at $z >2$ than at $z < 1$. +'onsidering the small probability of a given halo hosting a juasar at a given time. the more appropriate measure for justering of quasars is provided. by the unweighted halo correlation function.," Considering the small probability of a given halo hosting a quasar at a given time, the more appropriate measure for clustering of quasars is provided by the unweighted halo correlation function." + Recent estimates of the quasar correlation function. rough based on small datasets. show that clustering of quasars is indeed stronger at higher redshifts (Ixundic 1997: La Franea. Anclreant ancl Cristiani 1997).," Recent estimates of the quasar correlation function, though based on small datasets, show that clustering of quasars is indeed stronger at higher redshifts (Kundic 1997; La Franca, Andreani and Cristiani 1997)." + Authors of the second paper claim that the observed evolution of quasar clustering cannot be explained in all models of ACN activity., Authors of the second paper claim that the observed evolution of quasar clustering cannot be explained in all models of AGN activity. + Our analvsis suggests that this need not be true., Our analysis suggests that this need not be true. + Quasars must show decreasing clustering amplitude if the ACN activity is not correlated with the large scale environment., Quasars must show decreasing clustering amplitude if the AGN activity is not correlated with the large scale environment. + The assumption of a mass threshold also applies very well to groups and clusters of galaxies., The assumption of a mass threshold also applies very well to groups and clusters of galaxies. + Therefore. the cluster correlation. function. should. show a higher amplitucle at ueher recdshifts.," Therefore, the cluster correlation function should show a higher amplitude at higher redshifts." + We can use a cillerent ansatz: we can assume that the »ightest galaxies reside in halos that have formed relatively recently., We can use a different ansatz: we can assume that the brightest galaxies reside in halos that have formed relatively recently. + We can identify the mass scale Al.(2) by assuming hat these correspond to (Αν)cz1., We can identify the mass scale $M_\ast(z)$ by assuming that these correspond to $\nu(M_\ast) \simeq 1$. + For hierarchical nmocels. we know that a(Aly.2)7a(AlocMj. 2). implving hat v(Alo.z)woAl).," For hierarchical models, we know that $\sigma(M_1,z) > \sigma(M_2 > M_1,z)$ , implying that $\nu(M_2,z) > +\nu,M_1)$." +" As a high £6 corresponds to a veh correlation bias. it is clear that the halo correlation ""unction with threshold. Ad; will have a lower amplitude han the halo correlation function with threshold. A»."," As a high $\nu$ corresponds to a high correlation bias, it is clear that the halo correlation function with threshold $M_1$ will have a lower amplitude than the halo correlation function with threshold $M_2$." +" This indicates that the correlation function of Ad, halos at τι will have a much lower amplitude than its equivalent at το.", This indicates that the correlation function of $M_\ast$ halos at $z_1$ will have a much lower amplitude than its equivalent at $z_2$. + Therefore we expect rapid evolution of the galaxy correlation function. much faster than the linear growth rate. if galaxies correspond to A. halos.," Therefore we expect rapid evolution of the galaxy correlation function, much faster than the linear growth rate, if galaxies correspond to $M_\ast$ halos." + This may be a good model [or evolution of clustering of faint galaxies (Brainercl.SmailanclAloulcl 1995)., This may be a good model for evolution of clustering of faint galaxies \cite{evclus}. + The above discussion shows that the evolution of galaxy clustering is likely to depend: strongly on our choice of relation between halos and galaxies at different redshifts., The above discussion shows that the evolution of galaxy clustering is likely to depend strongly on our choice of relation between halos and galaxies at different redshifts. + Lf we identify similar objects at dillerent. redshifts. then these will show either a decreasing correlation function or nearly constant. clustering in comoving co-ordinates.," If we identify similar objects at different redshifts, then these will show either a decreasing correlation function or nearly constant clustering in comoving co-ordinates." + On the other hand. if we are working with the most prominent/numerous objects at cach redshift then we should see rapid evolution of clustering.," On the other hand, if we are working with the most prominent/numerous objects at each redshift then we should see rapid evolution of clustering." + Gastrophysical effects. in general. make the relation of halos and galaxies a little fuzzy. and this will result in an uncertain linear combination of the two types of evolution of clustering.," Gastrophysical effects, in general, make the relation of halos and galaxies a little fuzzy, and this will result in an uncertain linear combination of the two types of evolution of clustering." + Considering the uncertain relation of the relative evolution of galaxy/halo ancl mass clustering. the rate of erowth of the correlation function. at. large. scales should. not be interpreted. as the linear. growth rate of density perturbations.," Considering the uncertain relation of the relative evolution of galaxy/halo and mass clustering, the rate of growth of the correlation function at large scales should not be interpreted as the linear growth rate of density perturbations." + Direct determination of cosmological parameters by assuming the two rates to be identical can lead to wrong results., Direct determination of cosmological parameters by assuming the two rates to be identical can lead to wrong results. + Observations show that the Inter-Galactic Alecdium (LOAD) is ionised at the highest redshifts we can probe using known quasars ancl galaxies., Observations show that the Inter-Galactic Medium (IGM) is ionised at the highest redshifts we can probe using known quasars and galaxies. + Scenarios for reionisation of the LAL all in two basic categories: lonisation by quasars. which ive à low number density ancl hence the LGAL has a very xchy structure at early times.," Scenarios for reionisation of the IGM fall in two basic categories: Ionisation by quasars, which have a low number density and hence the IGM has a very patchy structure at early times." + lonisation bv proto-galaxies. να galaxies or star clusters that are distributed uniformly. eading to a quick and homogeneous reionisation.," Ionisation by proto-galaxies, dwarf galaxies or star clusters that are distributed uniformly, leading to a quick and homogeneous reionisation." + We have shown in this paper that halos. irrespective of heir mass. cluster very strongly at carly times.," We have shown in this paper that halos, irrespective of their mass, cluster very strongly at early times." + Therefore. he reionisation of the IGM will be patchy at a scale much arger than nας where pis the number density of the sources of ionisation.," Therefore, the reionisation of the IGM will be patchy at a scale much larger than $n^{-1/3}$, where $n$ is the number density of the sources of ionisation." + The non-uniformity of halos at. carly imes is illustrated in Πο.1., The non-uniformity of halos at early times is illustrated in fig.1. + Lt is thought that the earliest. clusters of stars formed when low mass halos ζω210AL. collapse after cooling ον HH» line cooling (Leemarketal.1997)., It is thought that the earliest clusters of stars formed when low mass halos $M_{halo} \simeq 10^6 M_\odot$ collapse after cooling by $_2$ line cooling \cite{molcool}. +. LW these clusters of stars were responsible for reionisation of IGM. then. as these wave avery large number density. η1/3zmLOkpey oper. the ionisation structure of the IGM will be fairly homogeneous.," If these clusters of stars were responsible for reionisation of IGM, then, as these have a very large number density, $n^{-1/3} \approx 10$ $_{proper}$, the ionisation structure of the IGM will be fairly homogeneous." + llowever. UV radiation from the first. clusters. loads: to dissociation of Ilo molecules and hence the Jeans mass increases by a considerable amount. (Hlaiman.ReesandLoeb 1997).," However, UV radiation from the first clusters leads to dissociation of $_2$ molecules and hence the Jeans mass increases by a considerable amount \cite{destrmol}." +". The second. generation of collapsed object are like dwarl galaxies. AM;c107""M. and these have a slightly lower number density. n4%=ODM peprepere"," The second generation of collapsed object are like dwarf galaxies $M_{halo} \simeq 10^{8-9} M_\odot$ and these have a slightly lower number density, $n^{-1/3} \approx 0.1$ $_{proper}$." + A visual comparison with Πο. suggests that some parts of the Universe will be at a much greater distance from the nearest source of ionising radiation., A visual comparison with fig.1 suggests that some parts of the Universe will be at a much greater distance from the nearest source of ionising radiation. + Therefore. one may expect oxdchiness at scales of Τρ μη. or about I0Mpcesouoe.," Therefore, one may expect patchiness at scales of $1$ $_{proper}$ , or about $10$ $_{comov}$." + The carly stages of reionisation from such sources may be observed. using the 2lem tomography type of observations. hough at a scale smaller than that suggested. for quasars (Macau.AleiksinandRees1997).," The early stages of reionisation from such sources may be observed using the $21$ cm tomography type of observations, though at a scale smaller than that suggested for quasars \cite{tomo}." +. On the other hand. if quasars are responsible for ionising the LGAL then the scale of patchiness will be much bigger than that expected from he number clensity.," On the other hand, if quasars are responsible for ionising the IGM then the scale of patchiness will be much bigger than that expected from the number density." + The relative distribution of halos and the underlying mass in Πο. suggests that halos form preferentially along ilaments/pancakes or their intersections., The relative distribution of halos and the underlying mass in fig.1 suggests that halos form preferentially along filaments/pancakes or their intersections. + Denser filaments ancl pancakes can “resist” being ionisecl much more ellicientlv than the under dense regions., Denser filaments and pancakes can “resist” being ionised much more efficiently than the under dense regions. + Photo-ionisation in the under-dense region will increase the Jeans mass ancl suppress collapse of low mass halos Cefstathiou at late times., Photo-ionisation in the under-dense region will increase the Jeans mass and suppress collapse of low mass halos \cite{supress} at late times. + This may be the reason lor the apparent lack of dwarf galaxies in voids., This may be the reason for the apparent lack of dwarf galaxies in voids. + This ellect willnot be as prominent in the filaments as the recombination time is smaller. and. ionisation proceeds much more slowly along over-dense regions thanit does in the undersdense regions.," This effect willnot be as prominent in the filaments as the recombination time is smaller, and, ionisation proceeds much more slowly along over-dense regions thanit does in the under-dense regions." +"Due to the collisionless nature of the plasma, pressure anisotropies (Pj#P,, where || and 1 are determined with respect to the magnetic field) can form during reconnection.","Due to the collisionless nature of the plasma, pressure anisotropies $P_\parallel \neq P_\perp$, where $\parallel$ and $\perp$ are determined with respect to the magnetic field) can form during reconnection." +" Fermi acceleration in contracting islands (Drakeetal.2006) and adiabatic cooling as B decreases due to conservation of the magnetic moment µοςv?/B« both drive Pj>P,."," Fermi acceleration in contracting islands \citep{Drake06a} and adiabatic cooling as $B$ decreases due to conservation of the magnetic moment $\mu \propto v_\perp^2/B \propto +P_\perp/B$ both drive $P_\parallel > P_\perp$." +" When an anisotropy is formed P,/withB Pj>P, the tension in bent magnetic fields weakens.", When an anisotropy is formed with $P_\parallel > P_\perp$ the tension in bent magnetic fields weakens. +" The fluid momentum equation with an anisotropy becomes where p is the mass density of the plasma, v is the bulk velocity, and B is the magnetic field."," The fluid momentum equation with an anisotropy becomes where $\rho$ is the mass density of the plasma, $\mathbf{v}$ is the bulk velocity, and $\mathbf{B}$ is the magnetic field." +" For 6)=B, the pressure equation reduces to the standard MHD equation.", For $\beta_\parallel = \beta_\perp$ the pressure equation reduces to the standard MHD equation. +" When fj>B, the tension force is reduced.", When $\beta_\parallel > \beta_\perp$ the tension force is reduced. + At high £this reduction of tension force is noticeable for even slight anisotropies in the pressure., At high $\beta$this reduction of tension force is noticeable for even slight anisotropies in the pressure. +" For —B, large enough, the tension force drops to zero, βor even becomes negative."," For $\beta_\parallel-\beta_\perp$ large enough, the tension force drops to zero, or even becomes negative." +" Since the tension of field lines acts as a restoring force for Alfvénn waves in standard MHD, the negative sign causes this oscillation to become an instability known as the firehose instability (Parker1958) for: 'This instability is fueled by the free energy contained in the pressure anisotropy."," Since the tension of field lines acts as a restoring force for Alfvénn waves in standard MHD, the negative sign causes this oscillation to become an instability known as the firehose instability \citep{Parker58} for: This instability is fueled by the free energy contained in the pressure anisotropy." +" The firehose instability causes magnetic field lines to kink, which eventually relieves the pressure anisotropy by causing scattering."," The firehose instability causes magnetic field lines to kink, which eventually relieves the pressure anisotropy by causing scattering." +" Alternatively, when 6)—f, is negative and large enough in magnitude, other instabilities can occur."," Alternatively, when $\beta_\parallel-\beta_\perp$ is negative and large enough in magnitude, other instabilities can occur." +" The mirror-mode instability and the ion cyclotron instability both occur when P,>Pj.", The mirror-mode instability and the ion cyclotron instability both occur when $P_\perp > P_\parallel$. +" For larger fj the mirror-mode becomes unstable at smaller values of [0—81| than the ion cyclotron mode, so the marginal mirror-mode criterion acts as the boundary between the stable and unstable regions."," For larger $\beta_\parallel$ the mirror-mode becomes unstable at smaller values of $|\beta_\parallel-\beta_\perp|$ than the ion cyclotron mode, so the marginal mirror-mode criterion acts as the boundary between the stable and unstable regions." +" Based on fluid theory assuming Τε=T; (Hasegawa1969),, the mirror-mode instability occurs when There are also kinetic modifications that can be made to the marginal instability criteria for firehose, mirror-mode, and ion cyclotron which make them more accurate."," Based on fluid theory assuming $T_e = T_i$ \citep{Hasegawa69}, the mirror-mode instability occurs when There are also kinetic modifications that can be made to the marginal instability criteria for firehose, mirror-mode, and ion cyclotron which make them more accurate." +" Although a rigorous analytic theory is not available, there are models that approximate the instability very well (Hellingeretal.2006;Bale"," Although a rigorous analytic theory is not available, there are models that approximate the instability very well \citep{Hellinger06, Bale09}." +" However, for simplicity we will just consider the conditions2009).. based on fluid theory."," However, for simplicity we will just consider the conditions based on fluid theory." +" In this study, we simulate several stacked current sheets similar to the compressed sectored heliospheric fields and associated current sheets, and follow the development of reconnection and islands."," In this study, we simulate several stacked current sheets similar to the compressed sectored heliospheric fields and associated current sheets, and follow the development of reconnection and islands." + We implement this system in a two-dimensional particle-in-cell (PIC) code and vary the temperature of the background plasma to test the dependence on 8., We implement this system in a two-dimensional particle-in-cell (PIC) code and vary the temperature of the background plasma to test the dependence on $\beta$. +" We observe that in finite Be systems (f> 0.5), very elongated islands form as opposed to the modest-aspect-ratio islands found at low Be (Be« where B, is the & based on the electron pressure."," We observe that in finite $\beta_e$ systems $\beta_e > 0.5$ ), very elongated islands form as opposed to the modest-aspect-ratio islands found at low $\beta_e$ $\beta_e < 0.5$ ), where $\beta_e$ is the $\beta$ based on the electron pressure." +" At 0.5),high 5, the increased due to the Fermi reflection of electrons within islands P,saturates the normal modest-aspect-ratio islands."," At high $\beta$, the increased $P_\parallel$ due to the Fermi reflection of electrons within islands saturates the normal modest-aspect-ratio islands." + Fermi reflection in highly elongated islands is less efficient because of the increased bounce time of the electrons so these islands are able to reach finite amplitude., Fermi reflection in highly elongated islands is less efficient because of the increased bounce time of the electrons so these islands are able to reach finite amplitude. +" At late time, however, even these elongated islands exhibit anisotropy instabilities, from Fermi reflection of both ions and electrons."," At late time, however, even these elongated islands exhibit anisotropy instabilities, from Fermi reflection of both ions and electrons." +" As a result, late-time magnetic islands remain highly elongated and do not become round as in the low-G regime."," As a result, late-time magnetic islands remain highly elongated and do not become round as in the $\beta$ regime." + This result has significant implications for the structure of islands that would be measured in the heliosheath., This result has significant implications for the structure of islands that would be measured in the heliosheath. +" Although 6, is rather moderate in the heliosheath, we find a mass ratio dependence suggesting long islands for a broad range of 5, in realistic mass ratios."," Although $\beta_e$ is rather moderate in the heliosheath, we find a mass ratio dependence suggesting long islands for a broad range of $\beta_e$ in realistic mass ratios." + A large £9 however may be necessary to sustain theelongation of these islands., A large $\beta$ however may be necessary to sustain theelongation of these islands. + Our simulations are performed with the PIC code p3d., Our simulations are performed with the PIC code p3d. + The initial conditions consist of eight Harris current sheets (Harris1962) where the magnetic pressure balances the plasma pressure., The initial conditions consist of eight Harris current sheets \citep{harris62} where the magnetic pressure balances the plasma pressure. +" Each Harris sheet consists of a tanh(y/wo) and sech?(y/wo) profile along the $ direction for an €-directed magnetic field and the density, respectively."," Each Harris sheet consists of a $\tanh(y/w_0)$ and $\sech^2(y/w_0)$ profile along the $\mathbf{\hat{y}}$ direction for an $\mathbf{\hat{x}}$ -directed magnetic field and the density, respectively." + The peak density of the Harris sheets is no., The peak density of the Harris sheets is $n_0$. +" In addition, there is à uniform background population that has a density of np=0.2no."," In addition, there is a uniform background population that has a density of $n_b = 0.2n_0$." +" These simulations are done in two dimensions so 0/0z=0, where Z is plane, parallel to the initial current."," These simulations are done in two dimensions so $\partial/\partial z = 0$, where $\mathbf{\hat{z}}$ is out-of-plane, parallel to the initial current." + The code uses normalized units., The code uses normalized units. + The timescale is normalized to the ion cyclotron time Q4., The timescale is normalized to the ion cyclotron time $\Omega_{\text{ci}}^{-1}$. + 'The distance scales arenormalized to the ion inertial length d;— C/Wp;., The distance scales arenormalized to the ion inertial length $d_i = c/\omega_{\text{pi}}$ . +" Thus, the velocity is normalized to the Alfvénn speed v4."," Thus, the velocity is normalized to the Alfvénn speed $v_A$." + The magnetic field is normalized to the asymptotic value of the reversed magnetic field Bo., The magnetic field is normalized to the asymptotic value of the reversed magnetic field $B_0$. + The density is normalized to no., The density is normalized to $n_0$. +" The pressure is normalized to Py=nom;vÀ,Βᾷ/Απ.", The pressure is normalized to $P_0 = n_0m_iv_A^2 = B_0^2/4\pi$. + The temperature is normalized to To=mivà., The temperature is normalized to $T_0 = m_iv_A^2$. + In order to vary the ϐ of these simulations we vary the temperature of the background population Tj., In order to vary the $\beta$ of these simulations we vary the temperature of the background population $T_b$. + This background temperature is the same for both ions and electrons., This background temperature is the same for both ions and electrons. +" The Harris equilibrium is used to balance the sharp change in the magnetic field strength across the current sheets, while the background represents the PUIs and has the greatest influence on late-time reconnection dynamics."," The Harris equilibrium is used to balance the sharp change in the magnetic field strength across the current sheets, while the background represents the PUIs and has the greatest influence on late-time reconnection dynamics." +" We performed simulations for 6=0.2, 1, 2, 3, and 4.8, where ϐ is based on the pressure in the asymptotic field with density ny."," We performed simulations for $\beta = 0.2$, $1$, $2$, $3$, and $4.8$, where $\beta$ is based on the pressure in the asymptotic field with density $n_b$." + Each simulation was advanced for a time of 1200! with a time resolution dt= , Each simulation was advanced for a time of $120\Omega_{\text{ci}}^{-1}$ with a time resolution $dt = 0.004\Omega_{\text{ci}}^{-1}$. +"The simulations are on a 204.8d;x102.4d; domain with a 0.0040;,'.grid scale resolution of A;=A, 0.05d;.", The simulations are on a $204.8d_i\times 102.4d_i$ domain with a grid scale resolution of $\Delta_x = \Delta_y = 0.05d_i$ . +" In order to complete such large runs, unless otherwise specified, we used 25 for the mass ratio of the ions to electrons."," In order to complete such large runs, unless otherwise specified, we used $25$ for the mass ratio of the ions to electrons." + This makes it easier to resolve small electron scales., This makes it easier to resolve small electron scales. +" In order to lessen the separation between the field and particle timescales, we set the ratio of the speed of light to the Alfvénn speed, c/ca, to 25 (in the heliosheath a more realistic value is near 6000)."," In order to lessen the separation between the field and particle timescales, we set the ratio of the speed of light to the Alfvénn speed, $c/c_A$ , to $25$ (in the heliosheath a more realistic value is near $6000$ )." + Reconnection is insensitive to the value of c/c4., Reconnection is insensitive to the value of $c/c_A$ . +" We start with a half-thickness for the current sheet wo= 0.5d;, so that collisionless reconnection can begin from particle noise."," We start with a half-thickness for the current sheet $w_0=0.5d_i$ , so that collisionless reconnection can begin from particle noise." +" The temperature in the Harris sheet is 0.2579for both ions and electrons, and there is no guide field."," The temperature in the Harris sheet is $0.25T_0$for both ions and electrons, and there is no guide field." + The largest ϐ we simulated was 4.8 since the electron, The largest $\beta$ we simulated was $4.8$ since the electron +Radial color profiles were obtained by subtracting two radial lummosity profiles from each other.,Radial color profiles were obtained by subtracting two radial luminosity profiles from each other. + The bottom panels of Figs., The bottom panels of Figs. + | show color profiles of the colorsB-R.. and{.," \ref{profiles} show color profiles of the colors, and." +. The errors in the color profiles were estimated by adding in quadrature the errors in the surface brightness profiles., The errors in the color profiles were estimated by adding in quadrature the errors in the surface brightness profiles. + Most galaxies show color gradients and become bluer with increasing radius (declining profiles)., Most galaxies show color gradients and become bluer with increasing radius (declining profiles). + It is believed that the dust content in disk LSB galaxies is low (dB95)., It is believed that the dust content in disk LSB galaxies is low (dB95). + Despite having no CO and metallicity information of the galaxies m our sample we probably can assume that some bulge LSB galaxies are, Despite having no CO and metallicity information of the galaxies in our sample we probably can assume that some bulge LSB galaxies are +required to match the overall age dispersion. are not a &ood fit near the observed distribution medians and thus result. in very low Ίντο probabilities.,required to match the overall age dispersion are not a good fit near the observed distribution medians and thus result in very low K-S probabilities. + To counter this. and provide a slightly more conservative (ic. larger) upper limit to the possible real age dispersion. we allowed the input dispersion of the apparent age distribution to vary from the observed values of cldex and 0.41 dex (see Table 1). linding that a smaller value of 0.3dcdex gave a much more probable model in both cases.," To counter this, and provide a slightly more conservative (i.e. larger) upper limit to the possible real age dispersion, we allowed the input dispersion of the apparent age distribution to vary from the observed values of dex and 0.41 dex (see Table 1), finding that a smaller value of dex gave a much more probable model in both cases." + This represented the core of the apparent age distribution quite well (see bottom. row of Fig. 5))., This represented the core of the apparent age distribution quite well (see bottom row of Fig. \ref{plotmodels}) ). +" The derived parameters (and limits) are in fact rather insensitive to this procedure. because of the very low weight that is attached to objects in the tails of the apparent age distribution bv IX-S ""he nature of these outliers and whether they olfer any support to the idea of a real age spread is discussed further in Section 5.3."," The derived parameters (and limits) are in fact rather insensitive to this procedure, because of the very low weight that is attached to objects in the tails of the apparent age distribution by K-S The nature of these outliers and whether they offer any support to the idea of a real age spread is discussed further in Section 5.3." + Confidence intervals on the model parameters. were estimated by renormalising the probability grid so that the sum over all possible parameter. combinations was unity., Confidence intervals on the model parameters were estimated by renormalising the probability grid so that the sum over all possible parameter combinations was unity. + Contours containing arbitrary [ractions of the probability distribution were calculated. from this eric., Contours containing arbitrary fractions of the probability distribution were calculated from this grid. + Integrating the erid over one or other of the parameter axes gave estimated confidence intervals in one parameter., Integrating the grid over one or other of the parameter axes gave estimated confidence intervals in one parameter. + The extent of the probability grids. were larger along the disc lifetime axis (typically up to MMyr) than displaved in Fig., The extent of the probability grids were larger along the disc lifetime axis (typically up to Myr) than displayed in Fig. + 5.r to ensure that all probability was accumulated., \ref{plotmodels} to ensure that all probability was accumulated. + The results of comparing the models with the three ONC data samples are tabulated in. Table 20 ancl illustrated in Figs., The results of comparing the models with the three ONC data samples are tabulated in Table \ref{modelresults} and illustrated in Figs. + 4. and 5., \ref{probgrids} and \ref{plotmodels}. + Fig., Fig. + 4. demonstrates to what extent the derived: model parameters are sensitive to each of the observational constraints using the example of the Spitzer sample.," \ref{probgrids} + demonstrates to what extent the derived model parameters are sensitive to each of the observational constraints using the example of the Spitzer sample." + It can be seen that the disc lifetime is very strongly constrained by the observed fraction of stars with. discs., It can be seen that the disc lifetime is very strongly constrained by the observed fraction of stars with discs. + The real age spread is stronely constrained. by the age distribution of stars that have lost their discs. whereas the age cistribution of stars with discs rules out. parameter space caturing large age spreads and short disc lifetimes.," The real age spread is strongly constrained by the age distribution of stars that have lost their discs, whereas the age distribution of stars with discs rules out parameter space featuring large age spreads and short disc lifetimes." + The outcome is a consistent interpretation from all hree samples. varving in statistical significance as expected rom the cdillerent data set sizes.," The outcome is a consistent interpretation from all three samples, varying in statistical significance as expected from the different data set sizes." + Figure 5. shows that the ack of any dilference in the observed. age distributions of he stars with and without clises constrains the real age spread to be much lower than the observed. age spread and ormallv consistent with zero for all three samples., Figure \ref{plotmodels} shows that the lack of any difference in the observed age distributions of the stars with and without discs constrains the real age spread to be much lower than the observed age spread and formally consistent with zero for all three samples. + The most constraining dataset is the large Spitzer sample. which demands that the real age spread be κ O.14ddex with 99 per cent confidence.," The most constraining dataset is the large Spitzer sample, which demands that the real age spread be $<0.14$ dex with 99 per cent confidence." + Even the smaller L-bancl sample provides a 99 per cent upper limit to the real age spread of «0.25 ddex., Even the smaller $L$ -band sample provides a 99 per cent upper limit to the real age spread of $<0.25$ dex. + The disc lifetime. as parameterised in this nmiocdel. is also well constrained in the two larger datasets at about 4d1 MMyr CA-band sample) or 6x1 MMyr(Spitzer sample). corresponding to a half-life of MMyr.," The disc lifetime, as parameterised in this model, is also well constrained in the two larger datasets at about $4\pm 1$ Myr $K$ -band sample) or $6\pm +1$ Myr(Spitzer sample), corresponding to a half-life of Myr." + H is larger for the L-band sample at 11+3 MMyr. due to the higher L-band disc frequency.," It is larger for the $L$ -band sample at $11\pm 3$ Myr, due to the higher $L$ -band disc frequency." + The sensitivity. of the conclusions {ο various moctel assumptions has been investigated., The sensitivity of the conclusions to various model assumptions has been investigated. + Ln particular. the possibilities of salvaging a real age spread that is anywhere near comparable with the observed age spread were explored in detail.," In particular, the possibilities of salvaging a real age spread that is anywhere near comparable with the observed age spread were explored in detail." + Figure 6 shows plots equivalent to those in Fig.5.. applied to the Spitzer sample. corresponding to the following model alterations.," Figure \ref{extras} shows plots equivalent to those in \ref{plotmodels}, , applied to the Spitzer sample, corresponding to the following model alterations." +the domain where coronal waves are observed. below 200 Maia (e.g.Warmthetal.2001a.2005:Virsuak is blocked by the coronagraph occultiug disk.,"the domain where coronal waves are observed, below 200 Mm \citep[e.g.][]{Warmuth04a, Warmuth05, Vrsnak05a, WillsDavey07, Patsourakos09a} is blocked by the coronagraph occulting disk." + Thus. a combination of EUVI and CORLL observations. coupled with nuneneal simulations. is required to develop a coherento picture of the carly stages of CME evolution in the low corona.," Thus, a combination of EUVI and COR1 observations, coupled with numerical simulations, is required to develop a coherent picture of the early stages of CME evolution in the low corona." + Iu this paper. we preseut an analysis of observations conibined with a elobal uuuerical simulation of a coronal wave event observed by STEREO on 13 February 2009.," In this paper, we present an analysis of observations combined with a global numerical simulation of a coronal wave event observed by STEREO on 13 February 2009." + This is the first nunerical simulation of au EUW coronal wave based on this observation in quadrature., This is the first numerical simulation of an EUV coronal wave based on this observation in quadrature. + We use a elobal MIID iodel for the solar corona that is driveu by real mmaeuctoeraim data. aud we drive the CME in a realistic way that matches the observations.," We use a global MHD model for the solar corona that is driven by real magnetogram data, and we drive the CME in a realistic way that matches the observations." + Numerical studies ofcoronal waves have been carried out previously., Numerical studies of coronal waves have been carried out previously. + However. these models were either two-dimensional (e.g.Pomoclletal. 2008).. considered onlv the expanding flux tube. ucelecting interaction with the surroundings (Delaunéeetal.2008).. sinulated the imteraction of coronal waves with only a local active region (Uchida1971:Ofiman&Thompson 2002).. or thev drove the coronal wave bv a pressure pulse (Wane2000:Wiretal. 2001).," However, these models were either two-dimensional \citep[e.g.][]{Pomoell08}, considered only the expanding flux tube, neglecting interaction with the surroundings \citep{Delannee08}, simulated the interaction of coronal waves with only a local active region \citep{Uchida74,Ofman02}, or they drove the coronal wave by a pressure pulse \citep{Wang00, Wu01}." +. By constraining the simmlatioun as much as possible with the observations. we obtain a full picture of the three-dimensioual evolution of the coronal magnetic field during the eruption. including those regions that are uot observed by white-lelt coronagraphis.," By constraining the simulation as much as possible with the observations, we obtain a full picture of the three-dimensional evolution of the coronal magnetic field during the eruption, including those regions that are not observed by white-light coronagraphs." + In this way. we hope to shed some light on the uature of coronal waves. their relationship to CATES. aud their theoretical description.," In this way, we hope to shed some light on the nature of coronal waves, their relationship to CMEs, and their theoretical description." + We describe the observations of the CME eveut iu Section 2.. and the uunuerical simulation iu Section 3..," We describe the observations of the CME event in Section \ref{sec:Obs}, and the numerical simulation in Section \ref{sec:Model}." + We prescut the results iu Section | aud discuss the iuplications for the various descriptions of coronal waves in Section 5.., We present the results in Section \ref{sec:Results} and discuss the implications for the various descriptions of coronal waves in Section \ref{sec:Discussion}. + We conclude our fiudiugs in Section 6.., We conclude our findings in Section \ref{sec:Conclusions}. + The coronal wave CAIE dinunines event on 13 February 2009 occurred when 9the STEREO spacecraft iad a separation angle of 917., The coronal wave – CME – dimmings event on 13 February 2009 occurred when the STEREO spacecraft had a separation angle of $^{\circ}$. + The Sun was in the very quictest part of its cycle. just at the start of the rise shase of solar excle 21.," The Sun was in the very quietest part of its cycle, just at the start of the rise phase of solar cycle 24." + NOAA active region 1012 was tle ouly active region on the solar disk., NOAA active region 1012 was the only active region on the solar disk. + Thus the majority of the surrounding environment was quiet Sun. with a ow-latitude coronal hole to the East of the active region (sce top panels Figure 1)).," Thus the majority of the surrounding environment was quiet Sun, with a low-latitude coronal hole to the East of the active region (see top panels Figure \ref{fig:f1}) )." +" Pre-cruption. the active region rosted a forward ""S sigmoid."," Pre-eruption, the active region hosted a forward “S” sigmoid." + The Sun produced à CATE hat was seen by spacecraft A right on the East limb. aud w spacecraft D. expaudiug from the ceuter of the disk.," The Sun produced a CME that was seen by spacecraft A right on the East limb, and by spacecraft B expanding from the center of the disk." + The CORLL iustrunents are internally occulted coronagraphs aud observe the ner solar corona in white heht from 1.3 - ον (Thompsonetal.2003)., The COR1 instruments are internally occulted coronagraphs and observe the inner solar corona in white light from 1.3 - 4 $_{\odot}$ \citep{Thompson03}. +.. Base difference mages (where a pre-event iniage at 05:15 UT. is subtracted from all subsequent images) of CORI-A data. are shown in the left panels of Figure 2...," Base difference images (where a pre-event image at 05:45 UT, is subtracted from all subsequent images) of COR1-A data, are shown in the left panels of Figure \ref{fig:f2}. ." +" The CORI images were taken with a temporal cadence of 10 uunutes. aud have a pixel size of 7.5""."," The COR1 images were taken with a temporal cadence of 10 minutes, and have a pixel size of $^{\prime\prime}$." + CORI-D also observed the CATE as a halo eveut. first becomine apparent bevond the occulting disk at 06:55 UT.," COR1-B also observed the CME as a halo event, first becoming apparent beyond the occulting disk at 06:55 UT." + The uou-differeuced CORI-A data (not shown) show a helinet streamer located north of AR 1012. and open streamers to the south.," The non-differenced COR1-A data (not shown) show a helmet streamer located north of AR 1012, and open streamers to the south." +" The ΕΙΝΤ tuagers observe the Sun out to 1.7 R.. aud produce «2018 nuages with a pixel size of 1.6"" (Wuclseretal.2001)."," The EUVI imagers observe the Sun out to 1.7 $_{\odot}$, and produce $\times$ 2048 images with a pixel size of $^{\prime\prime}$ \citep{Wuelser04}." +.. We analyze the 195 EEUVI images. which have a temporal cadeuce of 10 mnünuutes," We analyze the 195 EUVI images, which have a temporal cadence of 10 minutes." + Base difference images from EUVI-D iie shown iun Figure D., Base difference images from EUVI-B are shown in Figure \ref{fig:f3}. + The base difference nuages are produced bv first compensating for the solar rotation using SolarSoft' routine (hittp:/Awww.liusal.com/solarsoft}. so that all images are de-rotated to the pre-event nuage time at 05:15 UT.," The base difference images are produced by first compensating for the solar rotation using 's routine (http://www.lmsal.com/solarsoft), so that all images are de-rotated to the pre-event image time at 05:15 UT." + Then. the pre-event tuage is subtracted from all subsequeut dmaees.," Then, the pre-event image is subtracted from all subsequent images." + Base difference unaees highheht real intensity changes. with bright areas showing au merease in enussion with respect to the pre-event data.," Base difference images highlight real intensity changes, with bright areas showing an increase in emission with respect to the pre-event data." +" As well as showing the coronal wave bright front. base difference images also show depletious in iuteusity. known as ""coronal diuunuugs."," As well as showing the coronal wave bright front, base difference images also show depletions in intensity, known as “coronal dimmings”." + These are regions where the plasma deusitv has dramatically decreased due to plasma evacuation along the “opened” magnetic field. usually occumiug during an eruption (c.e.Iudsonοἳal.1996:Tarra&Sterling2001:ctal. 2007).," These are regions where the plasma density has dramatically decreased due to plasma evacuation along the “opened” magnetic field, usually occurring during an eruption \citep[e.g.][]{Hudson96, Harra01, Harra07}." +. Various works have shown that the core coronal dimming regious (i.c. black regions. bottom panels of Figure 1)). located on either side of the bright post-cruptive arcade (PEA) mark the footpoiuts of the crupting flux rope (e.g.WebbAttrilletal.2006:MeIutosh 2007).," Various works have shown that the core coronal dimming regions (i.e. black regions, bottom panels of Figure \ref{fig:f1}) ), located on either side of the bright post-eruptive arcade (PEA) mark the footpoints of the erupting flux rope \citep[e.g.][]{Webb00, Mandrini05, Crooker06, Attrill06, McIntosh07}." +. The inteusity drop iu core coronal dinuuiugs is substautial (typically[0-60%:ce.Chertok&Carechney 2005)., The intensity drop in core coronal dimmings is substantial \citep[typically $\sim$ 40 - 60\%; e.g.][]{Chertok05}. +. We process the EUVI base difference data using the automatic dinuuiues detection algorithuun described iu Attrill&Wil, We process the EUVI base difference data using the automatic dimmings detection algorithm described in \cite{Attrill09b}. +le-Davey(2009).. Figure E shows the output from this algorithin., Figure \ref{fig:f4} shows the output from this algorithm. + Iu addition to the core dinuuiugs near to the PEA. secondary. dinunines are also detected which develop remote from the active region and are spread across the solar disk.," In addition to the core dimmings near to the PEA, secondary dimmings are also detected which develop remote from the active region and are spread across the solar disk." + These secoudary dinuuiugs are nore subtle than the deep. core dimuuines aud are uot easv to identifv by eve in the base difference data.," These secondary dimmings are more subtle than the deep, core dimmings and are not easy to identify by eye in the base difference data." + We discuss the secoudary dininuings further in 50.1., We discuss the secondary dimmings further in \ref{subsec:sec_dims}. + The base difference images also show relatively concentrated. persistent brigltenines at the edge of the deep. core dinuuings.," The base difference images also show relatively concentrated, persistent brightenings at the edge of the deep, core dimmings." +" Examination of the location of coronal holes with respect to these briehteniugs shows that two brightenines (narked “A” aud ""D in Figure 1)) are situated along the boundary of a low-latitude coronal hole extending to the East of the active region.", Examination of the location of coronal holes with respect to these brightenings shows that two brightenings (marked “A” and “B” in Figure \ref{fig:f1}) ) are situated along the boundary of a low-latitude coronal hole extending to the East of the active region. +" Additionally. Irightenines ""D and 7€7 are located atthe edge of the core οο,"," Additionally, brightenings “B” and “C” are located atthe edge of the core dimmings." + Tn order to simulate the CME event. we use the Solar Corona (SC) model developed at the University of Michigau (Rousseyetal.2003b:Cohen2007. 2008b)..," In order to simulate the CME event, we use the Solar Corona (SC) model developed at the University of Michigan \citep{roussev03b,cohen07, cohen08b}. ." + The model i$ based on the global AMID, The model is based on the global MHD +cohunn 5 were arrived at using the values in col 7 of Table Las a enide aud requiring the 0 vs Age relation to be linear.,column 5 were arrived at using the values in col 7 of Table 1 as a guide and requiring the $\theta$ vs Age relation to be linear. + c) Rotation )oriod The resulting @ vs Age relation is shown iu Fie., c) Rotation period The resulting $\theta$ vs Age relation is shown in Fig. + 2., 2. + The slope of this line gives a rotation period for the mcleus of NGC LOGS of 1.8«τοῦ vr. which is a factor of 2.5 shorter than that found by Alloinetal.(2001) for the ceutral torus.," The slope of this line gives a rotation period for the nucleus of NGC 1068 of $1.8\times10^{6}$ yr, which is a factor of 3.3 shorter than that found by \citet{all01} for the central torus." + This value is assumed to be more accurate than that given carlicr (Bell2002a) since the earlier age calculations did not include accurate corrections for all ejection velocities., This value is assumed to be more accurate than that given earlier \citep{bel02a} since the earlier age calculations did not include accurate corrections for all ejection velocities. + Since the age is calculated from the preseut position of the objects and their ejection velocity perpendicular to the Lo-s. this would have affected the previous aee calculations.," Since the age is calculated from the present position of the objects and their ejection velocity perpendicular to the l-o-s, this would have affected the previous age calculations." + The intrinsic redshifts calculated in Table 3 are plotted versus level (triplet) iu Fie., The intrinsic redshifts calculated in Table 3 are plotted versus level (triplet) in Fig. + 3 where the redshift spacing between levels can be secn to increase. in qiailtiples of the redshift increment ΑΔ: = 0.062., 3 where the redshift spacing between levels can be seen to increase in multiples of the redshift increment $\Delta z$ = 0.062. +" Betore the Doppler compoucuts due to triplet ejection were removed here. the redshifts (zi, 44) were found to exhibit similar characteristics (Bell2002a).. except iu tha Case the redshift increments between evels were found to be inultiples of A: = 0.05,"," Before the Doppler components due to triplet ejection were removed here, the redshifts $_{\rm mean}$ ) were found to exhibit similar characteristics \citep{bel02a}, except in that case the redshift increments between levels were found to be multiples of $\Delta z$ = 0.05." + As long as the ejection velocities are σα]. are the rotation axis is tipped toward the observer. the slightly arecr increment of A: = 0.062 found rere after removal of the Doppler components would be anticipated.," As long as the ejection velocities are small, and the rotation axis is tipped toward the observer, the slightly larger increment of $\Delta z$ = 0.062 found here after removal of the Doppler components would be anticipated." + It is therofore not surprising that a similar relation was visible in the raw data., It is therefore not surprising that a similar relation was visible in the raw data. + This also reiuforces the likelihood that the results obtaiue here have not bee- introduced by the above caleulations., This also reinforces the likelihood that the results obtained here have not been introduced by the above calculations. + The uost interesting new result is hat when all Doppler components are accounted for. the intrinsic redshifts are all iiultiples ofthe redshift lucrenien Az = 0.062 first found by Burbidge over thirty vears ago (Durbidge1965:DurbideeaudDewitt 1990).," The most interesting new result is that when all Doppler components are accounted for, the intrinsic redshifts are all multiples of the redshift increment $\Delta$ z = 0.062 first found by Burbidge over thirty years ago \citep{bur68,bur90}." +. The intrinsic redshift values fouud for the four triplets near NGC 1068 are labeled A. D. C. aud D in Fie.," The intrinsic redshift values found for the four triplets near NGC 1068 are labeled A, B, C, and D in Fig." + 3., 3. + Two additional intrinsic redshift values (0.62 and O.31) can be predicted from Fig., Two additional intrinsic redshift values (0.62 and 0.31) can be predicted from Fig. + 3 corresponding to levels 1. aud 5., 3 corresponding to levels 4 and 5. + These appear to be the only other acceptable values since higher, These appear to be the only other acceptable values since higher +"very high values, so they span most of the interesting range for this quantity.","very high values, so they span most of the interesting range for this quantity." +" However, the validity of equation 1,, obtained in the case of gas in vertical hydrostatical equilibrium, must be checked in our simulations where energy is continually injected in discs."," However, the validity of equation \ref{eq:Pext}, obtained in the case of gas in vertical hydrostatical equilibrium, must be checked in our simulations where energy is continually injected in discs." + This is done in next sub-section., This is done in next sub-section. +" In figure 4 we report, for the four simulated galaxies, the relation between pressure and Σ]εοιά as found in the simulations."," In figure \ref{fig:P} we report, for the four simulated galaxies, the relation between pressure and $\Sigma_{\rm cold}$ as found in the simulations." +" We show in each panel the hydrodynamical pressure found in simulations, Psim, averaged in bins of the x-y grid as colored triangles, and the radial profile of the average of the same quantity as a line with a darker color."," We show in each panel the hydrodynamical pressure found in simulations, $P_{\rm sim}$, averaged in bins of the x-y grid as colored triangles, and the radial profile of the average of the same quantity as a line with a darker color." + The external pressure [οχι is computed using equation 1 from the radial profiles (colored dot-dashed lines).," The external pressure $P_{\rm + ext}$ is computed using equation \ref{eq:Pext} from the radial profiles (colored dot-dashed lines)." +" Moreover, to ease the comparison, each panel reports, as grey thick lines, the radial Xio34—Psim relations from the other simulations."," Moreover, to ease the comparison, each panel reports, as grey thick lines, the radial $\Sigma_{\rm + cold}-P_{\rm sim}$ relations from the other simulations." +" To obtain Pe, vertical velocity dispersions (v2) of cold gas and star particles have been computed in the galaxy frame."," To obtain $P_{\rm ext}$, vertical velocity dispersions $\langle v_z^2 \rangle$ of cold gas and star particles have been computed in the galaxy frame." +" For the stars, this quantity is equated to c2, while for the gas we use: (v2) +,"," For the stars, this quantity is equated to $\sigma_\star^2$, while for the gas we use: = v_z^2 + c_s^2 ," +color/magnitude euts would be considered one-halo pairs if the redshift were known.,color/magnitude cuts would be considered one-halo pairs if the redshift were known. + Furthermore. 19 pairs failed the color/magnitude cuts but passed the Az eut.," Furthermore, 19 pairs failed the color/magnitude cuts but passed the $\Delta z$ cut." + Therefore. the sample is complete alter the color/magnitude cuts aud contiuninated.," Therefore, the sample is complete after the color/magnitude cuts and contaminated." + For INC! collisions. we [found 2494 collisions passing the photometric color/magnitude cuts: 2020 of these passed the Az cuts. so would be assigned as pairs.," For INC collisions, we found 2494 collisions passing the photometric color/magnitude cuts; 2020 of these passed the $\Delta z$ cuts, so would be assigned as pairs." + 110.5 pairs passing the Az cuts failed the color magnitude cuts., 110.5 pairs passing the $\Delta z$ cuts failed the color magnitude cuts. + Therefore. this collision sample should be complete and Our approach is to estimate (he number of interlopers and apply the correction to (he munber of groups of 2 LRGs.," Therefore, this collision sample should be complete and Our approach is to estimate the number of interlopers and apply the correction to the number of groups of 2 LRGs." +" 568 objects from the photometric sample are <55” from a spectroscopic LRG object ancl pass the color/magnitude cuts at the spectroscopic redshift: 250 objects from the photometric sample are 55""<8Onoyr(2} and pass the color/magnitude cuts at the spectroscopic redshift.", 568 objects from the photometric sample are $\leq 55''$ from a spectroscopic LRG object and pass the color/magnitude cuts at the spectroscopic redshift; 250 objects from the photometric sample are $55'' < \theta < \theta_{max}(z)$ and pass the color/magnitude cuts at the spectroscopic redshift. + Based on the rates measured from spectroscopic collisions and after correcting for the expected number of interlopers NyepoiΝο. we expect. 709.3 of the remaining 789.1 collisions to be ‘true’ collisions (i.e.. would pass the Ac criterion as well. if the redshift were measured). and we expect to have missed 46.3 (rue collisions due to our color/maegnitude cuts.," Based on the rates measured from spectroscopic collisions and after correcting for the expected number of interlopers $N_{FB,fail} + N_{INC,fail}$, we expect 709.3 of the remaining 789.1 collisions to be `true' collisions (i.e., would pass the $\Delta z$ criterion as well, if the redshift were measured), and we expect to have missed 46.3 true collisions due to our color/magnitude cuts." + Therefore. we overestimate the total number of pairs by 62.4 (a Finally. we must estimate the total number of isolated LRGs (i.e.. Chose with no neighboring LRGs passing our CIC cuts) missing from our sample due to the incompleteness of the spectroscopic sample.," Therefore, we overestimate the total number of pairs by 62.4 (a Finally, we must estimate the total number of isolated LRGs (i.e., those with no neighboring LRGs passing our CiC cuts) missing from our sample due to the incompleteness of the spectroscopic sample." + Again making use of the assumption that objects in the photometric sample (hat would fail the color/magnitude cuts of our LRG sample are uncorrelated) with the spectroscopic sample and that the nunber of expected interlopers is negligible (Equs., Again making use of the assumption that objects in the photometric sample that would fail the color/magnitude cuts of our LRG sample are uncorrelated with the spectroscopic sample and that the number of expected interlopers is negligible (Eqns. + and 4)). we expect the total numberof isolated galaxies in the photometric sample to be Where Pyroup 15 the probability that an LRG is ina CiC group with Πω1.," \ref{failcontamFB} and \ref{failcontamINC}) ), we expect the total numberof isolated galaxies in the photometric sample to be where $p_{group}$ is the probability that an LRG is in a CiC group with $n_{sat} \geq 1$." + Using the observable οίκο: the number of objects from the photometric catalog not grouped with spectroscopic objects. along with the pass rate for a random set of LRG targets of36.64... allows us to solve for the ratio ροκ)αμ in the photometric sample.," Using the observable $N_{photo,iso}$, the number of objects from the photometric catalog not grouped with spectroscopic objects, along with the pass rate for a random set of LRG targets of, allows us to solve for the ratio $N_{pass}/N_{fail}$ in the photometric sample." +" Given the final estimate of the group multiplidty function (see below). we find Pyreup=0.132 and the isolated LRG contribution [rom the photometric sample. N,,,.*(1—Pyroup). is 2900."," Given the final estimate of the group multiplicity function (see below), we find $p_{group} = 0.132$ and the isolated LRG contribution from the photometric sample, $N_{pass}*(1-p_{group})$, is 2900." + This produces an increase in total objects in our sample of 856.. in agreement with the incompleteness rate reported by 2..," This produces an increase in total objects in our sample of , in agreement with the incompleteness rate reported by \citet{masjedi/etal:2006}." +each survey the one quoted in each paper. we found that the volume density of early-type galaxies is of the order of a few x I07Mpc7.,"each survey the one quoted in each paper, we found that the volume density of early-type galaxies is of the order of a few $\times$ $^{-5}$ $^{-3}$." + If we adopted the maximum redshift of each survey for which completeness 1s claimed. the number goes down to a few x 107*Mpe7?," If we adopted the maximum redshift of each survey for which completeness is claimed, the number goes down to a few $\times$ $^{-6}$ $^{-3}$." + We conclude from these numbers that it is perfectly plausible that all our isolated compact groups will end up as early-type galaxies in the field with no overpopulation Other indirect. evidence can further. strenghten our reasoning: likely candidates for such end-product from groups were previously observed by the MUSYC-YALE survey (van Dokkum et al..," We conclude from these numbers that it is perfectly plausible that all our isolated compact groups will end up as early-type galaxies in the field with no overpopulation Other indirect evidence can further strenghten our reasoning: likely candidates for such end-product from groups were previously observed by the MUSYC-YALE survey (van Dokkum et al.," + 2005)., 2005). + They are isolated galaxies. which. on deeper inspection. contain extended tidal tails and shells composed mainly of stars. with a small amount of residual gas. the so-called If we examine exisisting spectroscopic studies on galaxies in CGs. we find that observations of nearby isolated early-type galaxies (Collobert et al..," They are isolated galaxies, which, on deeper inspection, contain extended tidal tails and shells composed mainly of stars, with a small amount of residual gas, the so-called If we examine exisisting spectroscopic studies on galaxies in CGs, we find that observations of nearby isolated early-type galaxies (Collobert et al.," + 2006) have shown that the most massive galaxies in low density environments have abundance ratios similar to those of cluster galaxies., 2006) have shown that the most massive galaxies in low density environments have abundance ratios similar to those of cluster galaxies. + Mendes de Oliveira et al..," Mendes de Oliveira et al.," + 2005 also demonstrated that early type galaxies in HCGs are generally old., 2005 also demonstrated that early type galaxies in HCGs are generally old. + On the other hand. early-type galaxies in the field with small central velocity dispersions have properties that are consistent with extended episodes of star formation (Collobert et al..," On the other hand, early-type galaxies in the field with small central velocity dispersions have properties that are consistent with extended episodes of star formation (Collobert et al.," + 2006). as if coming from a past of multiple interaction and slow Available X-ray observations of CGs reveal a wide range of the X-ray diffuse emission. with a slight tendency for spiral-rich groups to have a small amount of X-ray emission (Ponman et al..," 2006), as if coming from a past of multiple interaction and slow Available X-ray observations of CGs reveal a wide range of the X-ray diffuse emission, with a slight tendency for spiral-rich groups to have a small amount of X-ray emission (Ponman et al.," + 1996)., 1996). + An opposite trend is tentatively detected for the HI content (Verdes-Montenegro et al..," An opposite trend is tentatively detected for the HI content (Verdes-Montenegro et al.," + 2001: Pompei et al..," 2001; Pompei et al.," + On the basis of these observations. the following scenario can be envisioned: galaxies belonging to more massive CGs evolve mainly within the group environment. giving rise to either a fossil group. 1e. a massive elliptical galaxy. surrounded by several dwarf galaxies. or to a field elliptical. whose abundance ratios are similar to. those observed in clusters.," On the basis of these observations, the following scenario can be envisioned: galaxies belonging to more massive CGs evolve mainly within the group environment, giving rise to either a fossil group, i.e. a massive elliptical galaxy, surrounded by several dwarf galaxies, or to a field elliptical, whose abundance ratios are similar to those observed in clusters." + As the galaxies have already evolved within the group. no or very' little young stellar population. should be present and they should continue to evolve viamergers.," As the galaxies have already evolved within the group, no or very little young stellar population should be present and they should continue to evolve via." + In both cases. the aforementioned X-ray emission i5 Galaxies belonging to less massive compact groups evolve in à more gradual way. probably by means of stripping of gas from each other through harassment. giving rise to several episodes of nuclear star formation. and diffuse HI emission within the group potential.," In both cases, the aforementioned X-ray emission is Galaxies belonging to less massive compact groups evolve in a more gradual way, probably by means of stripping of gas from each other through harassment, giving rise to several episodes of nuclear star formation, and diffuse HI emission within the group potential." + They will likely end up as a single isolated early-type galaxy with a younger stellar population, They will likely end up as a single isolated early-type galaxy with a younger stellar population +To calculate the structure aud evolution of an accreting flow. we solve (he equations of hydrodynamics where p is the mass density. /? is the gas pressure. v is the velocity. ο is the internal enerey density. and d is gravitational potential.,"To calculate the structure and evolution of an accreting flow, we solve the equations of hydrodynamics where $\rho$ is the mass density, $P$ is the gas pressure, ${\bf v}$ is the velocity, $e$ is the internal energy density, and $\Phi$ is gravitational potential." + We adopt an adiabatic equation of state P=(5-1l)e. where 5 is an aciabatic index.," We adopt an adiabatic equation of state $P~=~(\gamma-1)e$, where $\gamma$ is an adiabatic index." + Our calculations are performed in spherical polar coordinates (r.8.0).," Our calculations are performed in spherical polar coordinates $(r,\theta,\phi)$." + We assume axial symmetry aboutthe rotational axis of the accretion flow (6=0° and 130)., We assume axial symmetry aboutthe rotational axis of the accretion flow $\theta=0^\circ$ and $180^{\circ}$ ). + We perform simulations using (he pseucdo-Newtonian potential ® introduced by Paczvisski Wiita (1930) This potential approximates general relativistic effects in (he inner regions. lor a non- black hole.," We perform simulations using the pseudo-Newtonian potential $\Phi$ introduced by Paczyńsski Wiita (1980) This potential approximates general relativistic effects in the inner regions, for a non-rotating black hole." + In particular. the PaezvisskiWiita (P.W) potential reproduces (he last stable circular orbit at r=3424 as well as the marginally bound orbit al r=2R5.," In particular, the Paczyńsski–Wiita (P–W) potential reproduces the last stable circular orbit at $r=3 R_S$ as well as the marginally bound orbit at $r=2 R_S$." + Our standard computational domain isdefined to occupy the radial range r;=12Ry and the angular range 0°—0<1807., Our standard computational domain isdefined to occupy the radial range $r_i~=~1.5~R_S \leq r \leq \ r_o~=~ 1.2~R_B$ and the angular range $0^\circ \leq \theta \leq 180^\circ$. + We consider models with RG=107. Ro=10 Land RE=10.," We consider models with $R'_S=10^{-3}$, $R'_S=10^{-4}$, and $R'_S=10^{-5}$." + The r—0 domain is discretized into zones with 140. 180 and 220 zones (for RG=10.7. RG=10|. and RY=10? respectively) in the r direction and 200 zones in the @ direction.," The $r-\theta$ domain is discretized into zones with 140, 180 and 220 zones (for $R'_S=10^{-3}$, $R'_S=10^{-4}$, and $R'_S=10^{-5}$ respectively) in the $r$ direction and 200 zones in the $\theta$ direction." +" We fix zone size ratios. dryαι= 1.05. and d6;/d0,4=1.0 for OF«8 1807."," We fix zone size ratios, $dr_{k+1}/dr_{k}=1.05$ , and $d\theta_{l}/d\theta_{l+1} =1.0$ for $0^\circ \le \theta \le +180^\circ$ ." + For the initial conditions of the fIuid variables we follow PBO38a and adopt a Boncli, For the initial conditions of the fluid variables we follow PB03a and adopt a Bondi +estimates. mqo10747 cmand Line~LO’ Ix at rioBRS.,estimates $n_{\rm disc} \sim 10^{17}$ $^{-3}$ and $T_{\rm disc} \sim 10^5$ K at $r \sim 3R_S$. +" Applying these estimates to the magnetic pressure condition above. leads to a value for the field strength of a buovant ας tube within the inner regions of the disc of D,~5000 G. The observations provide two constraints on the densitv of the X-ray. emitting region."," Applying these estimates to the magnetic pressure condition above, leads to a value for the field strength of a buoyant flux tube within the inner regions of the disc of $B_c \sim 5000$ G. The observations provide two constraints on the density of the X-ray emitting region." + These arise from the constraints on the fitted column density and the ionisation parameter of the highly ionised absorber in PDS 456., These arise from the constraints on the fitted column density and the ionisation parameter of the highly ionised absorber in PDS 456. + ‘This is apparent in the hard. X-ray spectrum of PDS 456 above 7 keV. in the form of deep Ix-shell edges of highly ionised iron (ke and Fe XXV1)). present both in the current ELPIC spectrum (Reeves 2002. in preparation) and in earlier aand oobservations (Reeves 2000).," This is apparent in the hard X-ray spectrum of PDS 456 above 7 keV, in the form of deep K-shell edges of highly ionised iron (Fe and Fe ), present both in the current EPIC spectrum (Reeves 2002, in preparation) and in earlier and observations (Reeves 2000)." +" From the oobservations. the N-rav. spectrum of PDS 456 is well fit with a column density ng~vf, in the range 9107""Du - 5.lo""21 em 2implying; 9LO em: *«n«5.10T "," From the observations, the X-ray spectrum of PDS 456 is well fit with a column density $n_{\rm H} \sim n R_{\rm in}$ in the range $9\times 10^{23}$ - $5\times 10^{24}$ $^{-2}$ implying $9\times 10^8$ $^{-3}$ $< n <$ $5\times 10^9$ $^{-3}$." +"On the other hand. the ionisation parameter C2Lind; ies in the range 3101 - 2.5107. where La~LOM erg s is the hard. X-ray ionising Πακ Loading to the limits 410"" «nx 3.JIO em7 "," On the other hand, the ionisation parameter $U \simeq L_{\rm x}/n R^2_{\rm in}$ lies in the range $3\times 10^4$ - $2.5\times 10^5$, where $L_{\rm x} \sim 10^{45}$ erg $^{-1}$ is the hard X-ray ionising flux leading to the limits $4 \times 10^9$ $< n <$ $3 \times 10^{10}$ $^{-3}$." +"If we associate this density with the i00 post shock gas we obtain+ the estimateH n,5.«10Ln 7.", If we associate this density with the hot post shock gas we obtain the estimate $n_s \sim 5 \times 10^9$ $^{-3}$. + The evolution of the buovant [flux tubes provides a natural explanation for this very low value of the density of he X-rav emittinge region., The evolution of the buoyant flux tubes provides a natural explanation for this very low value of the density of the X-ray emitting region. +o As the {lux tubes emergeὃν from he disc and rise into the corona. gravitational downllow of plasma causes the density within the Dux tube to decrease o the point at which reconnection can take place efficiently.," As the flux tubes emerge from the disc and rise into the corona, gravitational downflow of plasma causes the density within the flux tube to decrease to the point at which reconnection can take place efficiently." + The density estimate above implies that the buovant [Lux ubes rise far into the corona before reconnection occurs., The density estimate above implies that the buoyant flux tubes rise far into the corona before reconnection occurs. + An estimate of the height at which reconnection occurs in erms of the pressure scale height of the disc (44) can be obtained from (Di Matteo 1998) In the case of a Shakura-Sunyaey accretion disc 1/2Ss implving that. fyRy., An estimate of the height at which reconnection occurs in terms of the pressure scale height of the disc $H$ ) can be obtained from (Di Matteo 1998) In the case of a Shakura-Sunyaev accretion disc $H/R \la 0.1$ implying that $h_{\rm flare} \ga R_{\rm in}$. +" Hence the total magnetic energy stored. in the corona of the inner disc should be of the order DZ5,(Sa~1073 erg.", Hence the total magnetic energy stored in the corona of the inner disc should be of the order $B_c^2 R_{\rm in}^3 / 8 \pi \sim 10^{51}$ erg. + This laree reservoir of stored accretion energy is enough to explain the large amplitude X-pav variations seen in PDS 456. and hishlights the elliciencv with which accretion energy must be converted to magnetic energy to explain these Daring events.," This large reservoir of stored accretion energy is enough to explain the large amplitude X-ray variations seen in PDS 456, and highlights the efficiency with which accretion energy must be converted to magnetic energy to explain these flaring events." + We now consider the rise time of the events in the licht. curves. (Aus30 ksec) in the light of the above model.," We now consider the rise time of the events in the light curves $t_{\rm rise} \ga +30$ ksec) in the light of the above model." + An attractive feature of the Petschek model is the short time scale on which reconnection can take place., An attractive feature of the Petschek model is the short time scale on which reconnection can take place. +" This can be of the order of a few Alfvénn times τν~ένRinldanmV)7AD, where / is the characteristic length scale associated with changes of the magnetic field (here /— t). ον ds the local Alfvénn speed ancl my, is the proton mass."," This can be of the order of a few Alfvénn times $\tau_{\rm A} \sim l / v_{\rm A} \sim R_{\rm in} +(4\pi n_s m_{\rm p})^{1/2}/ B_c$ where $l$ is the characteristic length scale associated with changes of the magnetic field (here $l \sim R_{\rm in}$ ), $v_{\rm A}$ is the local Alfvénn speed and $m_p$ is the proton mass." + ‘The estimates above viel τν~50 ksec. encouragingly close to the observed. rise time of the Hares.," The estimates above yield $\tau_{\rm A} \sim 50$ ksec, encouragingly close to the observed rise time of the flares." + The time scale τν provides a constraint on the density of the accretion dise and post shock gas., The time scale $\tau_{\rm A}$ provides a constraint on the density of the accretion disc and post shock gas. +" An accretion disc density significantly lower than raise~LOY or a shock density much higher than n,~QUU would increase τν to a level inconsistent with the observations.", An accretion disc density significantly lower than $n_{\rm disc} \sim 10^{17}$ $^{-3}$ or a shock density much higher than $n_s \sim 10^{10}$ $^{-3}$ would increase $\tau_{\rm A}$ to a level inconsistent with the observations. +" The energy. production rate for an individual Ilare can be estimated from (Di Matteo 1998): where L is the length. of the slow shock region. Zia is the N-ray temperature. and. m, is the electron mass."," The energy production rate for an individual flare can be estimated from (Di Matteo 1998):- where $L$ is the length of the slow shock region, $T_{\rm flare}$ is the X-ray temperature and $m_e$ is the electron mass." + The dimension L is the essentially the length of region of oppositelv-directed. magnetic field lines which we constrain to be of the order Lo—fy., The dimension $L$ is the essentially the length of region of oppositely-directed magnetic field lines which we constrain to be of the order $L \sim R_{\rm in}$. +" Adopting Zi~10"" ls we obtain Eus1075 erg s+."," Adopting $T_{\rm flare} \sim 10^9$ K we obtain $\dot{E}_{\rm flare} +\la 10^{43}$ erg $^{-1}$." + This suggests that the large scale X-ray variations in the light curve of PDS 456 involve @ 1000 or so individual Mares., This suggests that the large scale X-ray variations in the light curve of PDS 456 involve $\ga$ 1000 or so individual flares. + These Dares could not produce the large amplitude variations observed in PDS 456 if they were incoherent., These flares could not produce the large amplitude variations observed in PDS 456 if they were incoherent. + Llowever it is possible that the magnetic structure within the disc corona could reach a self-organized. critical state in which the reconnection and Uaring of one Ilux ube could. prompt similar fares in its neighbors. allowing a coherent cascade of Dares to develop (Leighly O'Brien 1991).," However it is possible that the magnetic structure within the disc corona could reach a self-organized, critical state in which the reconnection and flaring of one flux tube could prompt similar flares in its neighbors, allowing a coherent cascade of flares to develop (Leighly O'Brien 1997)." + The timescale on which the global cascade could take lace can be constrained from the Alfvénn time in the region involved. which is consistent with observed timescale of the laving events in the light curve.," The timescale on which the global cascade could take place can be constrained from the Alfvénn time in the region involved, which is consistent with observed timescale of the flaring events in the light curve." + This leads to à suggestion of why it is that PDS 456- releases such a large fraction of its accretion energy in the orm of laree-amplituce. coherent. X-ray variations.," This leads to a suggestion of why it is that PDS 456 releases such a large fraction of its accretion energy in the form of large-amplitude, coherent X-ray variations." + A Llare cascade along the lines of that suggested above would require hat magnetic energy was stored. in the disc corona until some trigeerine criterion was reached. promoting the first lave event.," A flare cascade along the lines of that suggested above would require that magnetic energy was stored in the disc corona until some triggering criterion was reached, promoting the first flare event." + Phe οποίον involved in such a cascade is likely o be governed by the time scale on which magnetic energy is pumped into the disc corona (μμ). and the time scale on which the triggering criterion is satisfied. (iis). , The energy involved in such a cascade is likely to be governed by the time scale on which magnetic energy is pumped into the disc corona $t_{\rm mag}$ ) and the time scale on which the triggering criterion is satisfied $t_{\rm trigger}$ +1997).,. +. Hybrid thermal/nonthermal models may also be a viable possibility for explaining the observed -ray spectrum., Hybrid thermal/nonthermal models may also be a viable possibility for explaining the observed $\gamma$ -ray spectrum. + The existeuce of such distributions is clearly established in the case of solar [ares ίο.ο..Coppi1999) aud it is therefore uatural to expect that similar distributious exist elsewhere in the universe.," The existence of such distributions is clearly established in the case of solar flares \citep[e.g.,][]{coppi99} and it is therefore natural to expect that similar distributions exist elsewhere in the universe." + Several such models have been discussecl in the literature 1999)..," Several such models have been discussed in the literature \citep[e.g.,][]{crider97,gierlinski97,poutanen96,poutanen98a,poutanen98b,coppi99}." + Fits to our broad-band spectrum using tle model of Poutanen&Sveussou(1996) provide a quantitative estinate of the electron distribution., Fits to our broad-band spectrum using the model of \citet{poutanen96} provide a quantitative estimate of the electron distribution. + In this case. the spectral data can be explained by a thermal Maxwellian distribution with au electron temperature of £T;=86 keV. coupled to a power-law electron spectrum with a power-law iudex of p.=L5.," In this case, the spectral data can be explained by a thermal Maxwellian distribution with an electron temperature of $kT_e = 86$ keV, coupled to a power-law electron spectrum with a power-law index of $p_e = 4.5$." +" The transition between the Maxwellian aud the power-law occurs at au electron. kinetic euergy of ~570 keV (75,45,=2.12).", The transition between the Maxwellian and the power-law occurs at an electron kinetic energy of $\sim570$ keV $\gamma_{min} = 2.12$ ). + The electron population is embedced in au accretion disk corona with au optical depth of 7=1.63., The electron population is embedded in an accretion disk corona with an optical depth of $\tau = 1.63$. + The derived spectral iudex for the power-law tail is somewhat harder (L5 versus 3.2) but still consistent. with that derived by Crideretal.(1997)., The derived spectral index for the power-law tail is somewhat harder (4.5 versus 3.2) but still consistent with that derived by \citet{crider97}. +. The role of z-productiou can also not be ruled out as a contribution to the nneasured spectrum., The role of $\pi^o$ -production can also not be ruled out as a contribution to the measured spectrum. + As initially proposed by Jordain&Roques(1991).. this ruechanism was used to explain the inadequacy of the STSO inodel at energies below 1 MeV. Π has not been applied to Cygnus at energies above 1. MeV. which would be required to test this moclel with the present. results.," As initially proposed by \citet{jourdain94}, this mechanism was used to explain the inadequacy of the ST80 model at energies below 1 MeV. It has not been applied to Cygnus X-1 at energies above 1 MeV, which would be required to test this model with the present results." + Models for advectiou-dominated accretion Hows (ADAF) predict ion temperatures of ~LOY ly in the inner part of the accretion [low (e.g..Chakrabarti&Titarchuk1995).. suggestiug the possibility of pion production (although it is not clear whether an ADAF is present during the low X-ray state ol Cyguusan X-1).," Models for advection-dominated accretion flows (ADAF) predict ion temperatures of $\sim10^{12}$ K in the inner part of the accretion flow \citep[e.g.,][]{chakrabarti95}, suggesting the possibility of pion production (although it is not clear whether an ADAF is present during the low X-ray state of Cygnus X-1)." + The spectrum presented here clearly iudicates the need for a non-Maxwellian electron euergy distribion., The spectrum presented here clearly indicates the need for a non-Maxwellian electron energy distribution. + In particular. it strougly suggests the presence ofa high euergy tail to that distribution.," In particular, it strongly suggests the presence of a high energy tail to that distribution." + Whether this results Grom thermal eeradients or from a more isothermal system coupled with a component remaius an open question., Whether this results from thermal gradients or from a more isothermal system coupled with a non-thermal component remains an open question. + The shape of the electrou clistribution aud its high euergv tail can only. be determined! by measurements that extencl into the MeV energy region., The shape of the electron distribution and its high energy tail can only be determined by measurements that extend into the MeV energy region. + There have been several studies of the broad-band hard X-ray emission from Cvenus S-1 based. in part. on OSSE data.," There have been several studies of the broad-band hard X-ray emission from Cygnus X-1 based, in part, on OSSE data." + These studies have concentrated on joint observatious with lower energy experiments. such as GINGA. ASCA or RXTE.," These studies have concentrated on joint observations with lower energy experiments, such as GINGA, ASCA or RXTE." + Since these observations are of relatively short curation. these spectra typically do uot exhibit a well-defined hard. tail due to limited statistics near ] MeV. Noneheless. the presence of a hard tail in the continuum emission uear 1 MeV now typears to be wel-estalished.," Since these observations are of relatively short duration, these spectra typically do not exhibit a well-defined hard tail due to limited statistics near 1 MeV. Nonetheless, the presence of a hard tail in the continuum emission near 1 MeV now appears to be well-established." + An excess above the standard Comptonization models has been discussed several times in the literature lor both Cygnus X-1 (e.g..MeConuelletal.1991:Philipsetal.1996:Linge1997). aud for GRO JO122+32 (vanDijketal.1995).," An excess above the standard Comptonization models has been discussed several times in the literature for both Cygnus X-1 \citep[e.g.,][]{mcconnell94,phlips96,ling97} and for GRO J0422+32 \citep{vandijk95}." +. The data presented here provide the best qlantitative rueasurement of the spectrum of Cyguus X-1 at these energies aud thus provide he best opportuuity to study the nature of this hard tail emission., The data presented here provide the best quantitative measurement of the spectrum of Cygnus X-1 at these energies and thus provide the best opportunity to study the nature of this hard tail emission. + The IBIS aud SPIiustrumeits οι INTEGRAL will provide only a mareinalfae) inproveiuent in the continuum, The IBIS and SPI instruments on INTEGRAL will provide only a marginal improvement in the continuum +The author thanks D. Sehmitt and D. Moss for providing electronic versions of their simulations. and Ix. Ferrierre. T. J. Jones. D. Clemens. and A. Pinnick for helpful comments and cliseussions.,"The author thanks D. Schmitt and D. Moss for providing electronic versions of their simulations, and K. Ferrièrre, T. J. Jones, D. Clemens, and A. Pinnick for helpful comments and discussions." + The author would also like to thank the anonvinous referee whose comments signilicantlv improved (his work., The author would also like to thank the anonymous referee whose comments significantly improved this work. + This work was partially supported by NSF grants AST 06-07500 and 09-07790 to Boston University. D. Clemens PI.," This work was partially supported by NSF grants AST 06-07500 and 09-07790 to Boston University, D. Clemens PI." +The dimensionless mass-acerction rate (or Eddiustou ratio) i=youAleτμ is defined in units of the Eddington luminosity ἔτι. where yeycOL is the niaxumuuni oeffücienev of απtephoton ΟΠΟΙΟΥ: conversion by the BI.,"The dimensionless mass-accretion rate (or Eddington ratio) $\dot{m}=\eta_{BH}\dot{M} c^2/L_{Edd}$ is defined in units of the Eddington luminosity $L_{Edd}$ , where $\eta_{BH} \simeq 0.1$ is the maximum efficiency of gravitational-to-photon energy conversion by the BH." + In the ADAF modelof (Navavan&Yi1995:Esinetal. 1998).. thermal radiation the accretion plasma at the level Ep=in?Lpggu//m.ds predicted when inHu~OL.," In the ADAF model \citep{ny95,esi98}, thermal radiation of the accretion plasma at the level $L_{T} = \dot{m}^2 L_{Edd}/\dot{m}_\ast$ is predicted when $\dot{m}\leq \dot{m}_\ast \sim 0.1$." +" The rate gpc1.5&107 Gs obtained if one formally equates Lp with the observed ZL,4,,4&LOores d."," The rate $\dot{m} \simeq 1.5 \times 10^{-5}$ is obtained if one formally equates $L_{T}$ with the observed $ L_{rad} +\simeq 10^{36}\,\rm erg\, s^{-1}$ ." + The compact radio eniüssiou cannot originate from au optically-thick accretion disk which would ouly be allowed or frou a much hotter ADAF plasina at Su mit is compatible with a svuchrotron origim (Liu.Petrosian.&Melia.2001).," The compact radio emission cannot originate from an optically-thick accretion disk which would only be allowed at $R \gg R_{rad} $, or from a much hotter ADAF plasma at smaller scales, but is compatible with a synchrotron origin \citep{lpm04}." +. The thermal output of the ADAF would be peaked mostly iu the hard X-ray/soft eanuna-rav domain. where the quiescent luninosity is at least 2 orders of mmaguitucde simaller than μι," The thermal output of the ADAF would be peaked mostly in the hard X-ray/soft gamma-ray domain, where the quiescent luminosity is at least 2 orders of magnitude smaller than $L_{rad}$." + This results i the accretion rate estimate ij4;1ο95x10Ὁ consistent with and observations., This results in the accretion rate estimate $\dmn = \dm /10^{-6} \lesssim 10^{-6} $ consistent with ADAF model and observations. + Even for ibi this small. ADAPintoLOMcu roni Ser ÀÁ.," It becomes the main contributor to the $\lesssim 100 \,\rm GeV$ Compton flux from the plerion at $R\gtrsim 10^{17}\,\rm cm$ from Sgr $^\ast$." + Synehrotron radiation from these sanie imulti-TeV electrons at δι10!αμ produces the quiesceut. X-rav flux from Ser ," Synchrotron radiation from these same multi-TeV electrons at $R\leq 10^{17}\,\rm cm$ produces the quiescent X-ray flux from Sgr $^\ast$." +Chandra observatious (Baganoffetal.2003) reveal diffuse thermal N-ravs with luuinosity 2.1«10?!eres1 from the central parsec region of Ser produced by AEzcL.3keV plasma with s.=26η)σωμα3 clensity. where gj;=y1/101 is the volun filling factor.," Chandra observations \citep{Chandra} reveal diffuse thermal X-rays with luminosity $2.4\times 10^{34}\,\rm erg \, s^{-1}$ from the central parsec region of Sgr $^\ast$, produced by $k T \simeq +1.3\,\rm keV$ plasma with $n_e =26/\eta_{f}^{1/2}\,\rm cm^{-3}$ density, where $\eta_f =\eta_{-1}/10^{-1} $ is the volume filling factor." + Equating the eas pressure vnAt in that region with the energv density of the wind gives the shock distance The TeV Luuimositv L.c107cres|o naplies a total electron. acceleration power L.>loeres t.," Equating the gas pressure $n_e kT$ in that region with the energy density of the wind gives the shock distance The TeV luminosity $L_\gamma \simeq 10^{35} \,\rm erg \,s^{-1}$ implies a total electron acceleration power $L_e \gtrsim 10^{36}\,\rm +erg \, s^{-1}$ ." + Asstuuine that the cficiency of electron acceleration LfLojanontlcimad&30%. aud letting the rest go to thermal plasiua and protons(unlike iu pulsar winds where relativistic electrons dominate). we arrive at L5;Z0.3," Assuming that the efficiency of electron acceleration $L_e /L_{wind} \lesssim 30\%$, and letting the rest go to thermal plasma and nonthermal protons (unlike in pulsar winds where relativistic electrons dominate), we arrive at $L_{w.37} \gtrsim +0.3\,$." + Note that Eq.(3) predicts that the maxima power of the outflow sustainable at clistances Raju.$ 1, CO formation is favoured against TiO. The presence of strong CN bands makes it difficult to use simple visual means to select stars with enhanced lithium, and even harder to measure the strength of the line." + Phe method chosen here to identifv the lithium line ina particular star was to subtract the mean observed spectrum (described above) of the spectral group appropriate for the star in question. scaled: according to the strength of the CN bands in the individual spectrum.," The method chosen here to identify the lithium line in a particular star was to subtract the mean observed spectrum (described above) of the spectral group appropriate for the star in question, scaled according to the strength of the CN bands in the individual spectrum." + The scaling factor was determined by minimizing the dillerence between the spectrum. and the scaled. mean., The scaling factor was determined by minimizing the difference between the spectrum and the scaled mean. + Lhe subtracted spectrum is mainly that of CN. but. includes. contributions from metal lines and other molecular bands.," The subtracted spectrum is mainly that of CN, but includes contributions from metal lines and other molecular bands." + Lhe subtraction process is designed to reveal only those features that deviate significantly from the scaled mean spectrum., The subtraction process is designed to reveal only those features that deviate significantly from the scaled mean spectrum. + Fig., Fig. + 2 shows examples of this process For three stars: two typical IN stars (Group 1). one with and the other without a lithium line.," \ref{montage} shows examples of this process for three stars: two typical N stars (Group 1), one with and the other without a lithium line," +The implications of the above results with regards to spacecraft navigational needs are twofold.,The implications of the above results with regards to spacecraft navigational needs are twofold. + First. there are nunierous sources available at ligh frequencies. aud these sources persist over relatively long (five vears) periods of time.," First, there are numerous sources available at high frequencies, and these sources persist over relatively long (five years) periods of time." + Second. the sources are more compact at these higher frequencies aud should provide high-quality astrometric reference points for spacecraft navigation.," Second, the sources are more compact at these higher frequencies and should provide high-quality astrometric reference points for spacecraft navigation." + Such sources also provide a foundation upon which to build a future CRE aud are of benefit to radio astrononiv. im general. as fiducial references.," Such sources also provide a foundation upon which to build a future CRF and are of benefit to radio astronomy, in general, as fiducial references." + Future nuage observatious should continue to improve our understauidiug of source structure at EK and Q baud and allow us to better characterize the variability of the structure over fiue and its effect on astrometric positional stability., Future imaging observations should continue to improve our understanding of source structure at K and Q band and allow us to better characterize the variability of the structure over time and its effect on astrometric positional stability. + This research was partially supported through NASA contracts with the California Institute of Techuologv and the U.S. Naval Observatory., This research was partially supported through NASA contracts with the California Institute of Technology and the U.S. Naval Observatory. + Tlis research macde use of the USNO Radio Reference. Frame Tage Database (RRFID)., This research made use of the USNO Radio Reference Frame Image Database (RRFID). +efficiency include a laree fraction of Pop III stars being verv massive stus (M>300M.) which may collapse to black holes instead of exploding as supernovae (Iakavy. Shaviv. Zinamon 1967; Bond. Arnett. Carr 1984: Glatzel. Fricke. El Eid L985; Woosley 1936) or not all ejected metals are able to be transported to IGM.,"efficiency include a large fraction of Pop III stars being very massive stars $M\ge 300\msun$ ) which may collapse to black holes instead of exploding as supernovae (Rakavy, Shaviv, Zinamon 1967; Bond, Arnett, Carr 1984; Glatzel, Fricke, El Eid 1985; Woosley 1986) or not all ejected metals are able to be transported to IGM." + We note Chat Model 4211 is close to the model used in Cen (2003). having perhaps more conventional (thus thought to be more reasonable) values for the relevant parameters.," We note that Model 1 is close to the model used in Cen (2003), having perhaps more conventional (thus thought to be more reasonable) values for the relevant parameters." + We also (test a variant of Model 11 by using WL)x(1+z2)/? instead of being constant with redshift and find τι=0.08 for that case versus 0.07 for a constant. V(17) (Model #11)., We also test a variant of Model 1 by using $\Psi(II)\propto (1+z)^{1/2}$ instead of being constant with redshift and find $\tau_e=0.08$ for that case versus $0.07$ for a constant $\Psi(II)$ (Model 1). + From Table 1 we can craw several conclusions., From Table 1 we can draw several conclusions. + First. without Pop III stars (Models #11.2). it is likely that 7.<0.09 for both LCDAIAP and LCDAIAP+ models. inconsistent with WALAP results al >2o level (Ixogut 2003).," First, without Pop III stars (Models 1,2), it is likely that $\tau_e < 0.09$ for both LCDMAP and LCDMAP+ models, inconsistent with WMAP results at $\ge 2\sigma$ level (Kogut 2003)." + The conclusions reached here with regard to the requirement of a top-heavv. IAIF lor Pop HI stars in order to explain the observed high 7. value are consistent with those by Wvithe Loeb (2003b). Haiman llolder (2003) and Sokasian (2003).," The conclusions reached here with regard to the requirement of a top-heavy IMF for Pop III stars in order to explain the observed high $\tau_e$ value are consistent with those by Wyithe Loeb (2003b), Haiman Holder (2003) and Sokasian (2003)." + Second. lor reasonable ranges for the star lormation efficiency [e(LL.HI).<0.1. οΠο.HE)< 0.002] and ionizing photon escape fraction CQLII)< 0.20) from Pop HI galaxies. we expect that 7.<0.12 for LCDAIAP mocdel (Model 4633) and τx0.15 for LCDMAP-- model (Model #44).," Second, for reasonable ranges for the star formation efficiency $c_*(HI,III)\le 0.1$, $c_*(H_2,III)\le 0.002$ ] and ionizing photon escape fraction $f_{esc}(III)\le 0.20$ ) from Pop III galaxies, we expect that $\tau_e \le 0.12$ for LCDMAP model (Model 3) and $\tau_e \le 0.15$ for LCDMAP+ model (Model 4)." +" Third. prolonging the Pop Η1 era bx making (he metal enrichment of the IGM from Pop III galaxies less efficient increases (he Thomson optical depth incrementally,"," Third, prolonging the Pop III era by making the metal enrichment of the IGM from Pop III galaxies less efficient increases the Thomson optical depth incrementally." + Finally. while 7.=0.17 is the mean value determined by WALAP observations. it seems hard to reach in the LCDAIAP model.," Finally, while $\tau_e=0.17$ is the mean value determined by WMAP observations, it seems hard to reach in the LCDMAP model." + There are two possible wavs to achieve a Thomson optical depth as high as 7.>0.17: either the ionizing photon escape Traction is lavee (fi(L4)> 0.3) and the underlving cosmological model has a blue power spectrum will a posilive Ull ton>1.03 (Model 455) orPop III star formation efficiency in minihalos is substantially larger than current simulations seem {ο indicate. requiring ον1ο.111)>0.01.," There are two possible ways to achieve a Thomson optical depth as high as $\tau_e\ge 0.17$: either the ionizing photon escape fraction is large $f_{esc}(III)\ge 0.3$ ) and the underlying cosmological model has a blue power spectrum with a positive tilt to $n\ge 1.03$ (Model 5) orPop III star formation efficiency in minihalos is substantially larger than current simulations seem to indicate, requiring $c_*(H_2,III)>0.01$." + Figures 1.2 show the reionization history aud cumulative Thomson optical depth. respectively. for four representative models listed in Table 1.," Figures 1,2 show the reionization history and cumulative Thomson optical depth, respectively, for four representative models listed in Table 1." + We see that without Pop ILLE stars (Mocel #11) relonization event only occurs once al ze6 and (he reionization history is monotonic., We see that without Pop III stars (Model 1) reionization event only occurs once at $z\sim 6$ and the reionization history is monotonic. + With Pop HI stars the first reionization occurs al 2=15—20 (Models. z:33.4).," With Pop III stars the first reionization occurs at $z=15-20$ (Models 3,4)." + In the ast case (Model 455) where Pop III star formation efficiency. in minihalos is substantially arger (han suggested by current. simulations of Abel. Brvan Norman (2002). 0.01 versus e»0.002. the first reionization occurs al 2e20—25.," In the last case (Model 5) where Pop III star formation efficiency in minihalos is substantially larger than suggested by current simulations of Abel, Bryan Norman (2002), $0.01$ versus $\sim 0.002$, the first reionization occurs at $z\sim 20-25$." +" While both Model 444 and Model 4655 wave Comparable total Thomson optical depth. 7,=0.15—0.16. their respective reionization stories al early (mes (2> 15) are somewhat different."," While both Model 4 and Model 5 have comparable total Thomson optical depth, $\tau_e=0.15-0.16$, their respective reionization histories at early times $z\ge 15$ ) are somewhat different." + We [ind (hat lor Models. #44. ninihalos and large halos contribute to and of all ionizing photons to reionize the universe al the [inst tme. by which 9 ionizine photons per barvon have been produced: the corresponding numbers for Model 4:55 are 61%... and 11 photons per barvon.," We find that for Models 4, minihalos and large halos contribute to and of all ionizing photons to reionize the universe at the first time, by which 9 ionizing photons per baryon have been produced; the corresponding numbers for Model 5 are , and 11 photons per baryon." + In contrast.," In contrast," +"where The induced dipole moment. Ding. is then given by an integral over the surface of the sphere: eiving We set ff,=ChH(w.m)functiMbOyexp( ikl). where C's(w.) is a normalizing lion. a=dV fd6. and f(r) satisfies with s=we|dei).","where The induced dipole moment, $D_{\mathrm{ind}}$, is then given by an integral over the surface of the sphere: giving We set $H_\phi = C_5 (\omega,\overline{\sigma}) N_1 (\theta) h(r) +\exp (-{\mathrm i} \omega t)$ , where $C_5 (\omega,\overline{\sigma})$ is a normalizing function, $N_1 = {\mathrm d} P_1 /{\mathrm d} \theta$ , and $h(r)$ satisfies with $\kappa^2 = \omega (\omega + 4\pi {\mathrm i} \overline{\sigma})$." + We choose the solution that represcnits Culgolng Waves. so llere hi((wr) is the spherical Bessel function defined in ? ," We choose the solution that represents outgoing waves, so Here $h_1^{(1)} (\kappa r)$ is the spherical Bessel function defined in \citet{absteg}. ." +The Maxwell equation V:H=ile|απισ)Ε thenelves For small r this becomes: Sut also. for small kr. we have CX2)). so In this appendix we have usec ordinary polar coordinates in [at space.," The Maxwell equation $\nabla\times {\mathbf H} = +-{\mathrm i} (\omega + 4\pi {\mathrm i} \overline{\sigma}){\mathbf E}$ thengives For small $r$ this becomes: But also, for small $r$, we have \ref{eq:E_r_B}) ), so In this appendix we have used ordinary polar coordinates in flat space." +" We need now to transform to coordinates of a ET fat space. with metric To eet from £L, of CX13)) to fis of (13)). in the Lat ERN. metric 6X15)). we first refer to 2.. section. 4.8. and multiply bv X."," We need now to transform to coordinates of a FRW flat space, with metric To get from $H_{\phi}$ of \ref{eq:H_B}) ) to $F_{12}$ of \ref{eq:F12def}) ), in the flat FRW metric \ref{eq:ds4}) ), we first refer to \citet{wein2}, section 4.8, and multiply by $\chi$ ." + We then convert from / to a and rto x by multiplying bv I7(p)., We then convert from $t$ to $\eta$ and $r$ to $\chi$ by multiplying by $R^2 (\eta)$. + We also convert gton. & to k and σ lom—finm We can obtain the corresponding function in curved space by multiplying (67)) by INí(0)exp(.μη): We can now [ind €'5(n.o) bv matching. CX16)). with CXI1T)) for small \: With P(r) given by (10)). this reduces to 1n the limit a-0. (AL9)) tends to (32)). as it should.," We also convert $\omega$ to$n$, $\kappa$ to $k$ and $\overline{\sigma}$ to $\sigma = R(\eta) \overline{\sigma}$: We can obtain the corresponding function in curved space by multiplying \ref{eq:f3out2perm}) ) by $N_1 (\theta) \exp(-{\mathrm i} n \eta)$: We can now find $C_3 (n,\sigma)$ by matching \ref{eq:F12flat}) ) with \ref{eq:F12curv}) ) for small $\chi$: With $D (n)$ given by \ref{eq:dipstrength}) ), this reduces to In the limit $\sigma \rightarrow 0$, \ref{eq:normC3perm2}) ) tends to \ref{eq:normC1}) ), as it should." + li is well known that although Alaxwell’s equations. are conformally invariant in d-dimensional space. the Lorenz condition is not (?.chapter2)..," It is well known that although Maxwell's equations are conformally invariant in 4-dimensional space, the Lorenz condition is not \citep[][chapter 2]{hoyle}." + So a theory thataspires tobe completely conformally invariant. such as the theoryof Mannheim discussedin section 16.. must include à moclified gauge condition if it is to describe the propagationof potentials.," So a theory thataspires tobe completely conformally invariant, such as the theoryof Mannheim discussedin section \ref{sec:mannheim_model},, must include a modified gauge condition if it is to describe the propagationof potentials." +" We note that the conformal weights of ο. να. Pand wl, are 2.4. Lane zero. respectively."," We note that the conformal weights of $g^{\mu \nu}$ , $\sqrt{g}$ , $\Phi$ and $A_{\nu}$ are $-2$ , $4$ ,$-1$ and zero, respectively." + Phe equation (44)), The equation \ref{eq:lrnz2}) + Phe equation (44))., The equation \ref{eq:lrnz2}) +ere only the Legendre polvnommals. Pj} and P2(Qi) depend on the direction k.,"Here only the Legendre polynomials, $P_r(\mu)$ and $P_{r'}(\mu')$ depend on the direction $\hat{\bf k}$." + To perform the integration dQ). we use the addition theorem for the Ισ GO)(CAT)," To perform the integration $d\Om_{\hat{\bf k}}$, we use the addition theorem for the spherical harmonics $Y_{rs}$, )." +" The orthogonality of the spherical harmonics then yields νο... n"").", The orthogonality of the spherical harmonics then yields ') = '). + Iu Eq. CÀA16)), In Eq. \ref{dT3}) ) +" the integration over do, then leads to terms of the form (n:n/P.(na’) and innP,n/).", the integration over $d\Om_{\hat{\bf k}}$ then leads to terms of the form $({\bf n}\cd{\bf n}')P_r({\bf n}\cd{\bf n}')$ and $({\bf n}\cd{\bf n}')^2P_r({\bf n}\cd{\bf n}')$. + To reduce them. we use d," To reduce them, we use ." +oppi Applving this aud its iteration for (A20)(7.(0). we obtain tete ander | AyE BelidkttyAUR)19) | joC£)0) (ety#9) ο d LUM UsA122) where the areuneut of the Leseudre polyuoiials. nay. has beeu suppressed.," Applying this and its iteration for $x^2 P_r(x)$, we obtain k^2dk dtdt' + (k(t_0-t')) - P_r [j_r(k(t_0-t)j_r”(k(t_0-t')) + j_r(k(t_0-t'))j_r”(k(t_0-t)) ”(k(t_0-t'))] -4 + ] '(k(t_0-t')), where the argument of the Legendre polynomials, $\bf n\cd n'$, has been suppressed." +" Usine the relations10À2 forBessel 3rfuuctious. Sleaud its iteration for 7"". we cau rewrite Eq. (A23))"," Using the relations forBessel functions, and its iteration for $j''$ , we can rewrite Eq. \ref{dT4}) )" +" iu terms of theBessel fuuctious j,2to j,|».", in terms of theBessel functions $j_{r-2}$to $j_{r+2}$ . +"At a distance of 16.7 Mpc (Meietal.2007), M87 is the dominant galaxy of the Virgo Cluster.","At a distance of 16.7 Mpc \citep{2007ApJ...655..144M}, M87 is the dominant galaxy of the Virgo Cluster." + It is one of the nearest radio galaxies and was the first extragalactic X-ray source to be identified., It is one of the nearest radio galaxies and was the first extragalactic X-ray source to be identified. +" Because of its proximity, many interesting astrophysical phenomena can be studied in more detail in M87 than in other comparable objects (see citetL999LNP...530....R. for an overview)."," Because of its proximity, many interesting astrophysical phenomena can be studied in more detail in M87 than in other comparable objects (see \\citet{1999LNP...530.....R} for an overview)." +" Among its many remarkable features is the several billion solar mass supermassive black hole at its centre (Macchettoetal.1997;Gebhardt&Thomas2009) and the prominent jet extending from the nucleus, visible throughout the electromagnetic spectrum."," Among its many remarkable features is the several billion solar mass supermassive black hole at its centre \citep{1997ApJ...489..579M, 2009ApJ...700.1690G} and the prominent jet extending from the nucleus, visible throughout the electromagnetic spectrum." +" The central regions of M87, in particular the structure of the jet, have been studied and compared intensively at radio, optical, and X-ray wavelengths 2010).."," The central regions of M87, in particular the structure of the jet, have been studied and compared intensively at radio, optical, and X-ray wavelengths \citep[e.g.][]{1991AJ....101.1632B, + 1996A+A...307...61M, 2001A+A...365L.181B, 2001ApJ...551..206P, + 2004ApJ...607..294S, 2005ApJ...627..140P, 2007ApJ...668L..27K, + 2008A+A...482...97S, 2010arXiv1003.5334W}." +" Compared to the available information at these wavelengths, our knowledge of M87 at far-infrared (FIR) wavelengths is rather poor."," Compared to the available information at these wavelengths, our knowledge of M87 at far-infrared (FIR) wavelengths is rather poor." +" A controversial issue is the origin of the FIR emission in M87, i.., the question of whether the FIR emission is caused entirely by synchrotron emission or whether there is an additional contribution from dust associated with either the global interstellar medium or a nuclear dust component."," A controversial issue is the origin of the FIR emission in M87, i.e., the question of whether the FIR emission is caused entirely by synchrotron emission or whether there is an additional contribution from dust associated with either the global interstellar medium or a nuclear dust component." + This question is partly driven by the observation of faint dust features in deep optical images (Sparksetal.1993;Ferrarese2006)..," This question is partly driven by the observation of faint dust features in deep optical images \citep{1993ApJ...413..531S, 2006ApJS..164..334F}." + Several papers on the FIR emission of M87 arrive at different conclusions., Several papers on the FIR emission of M87 arrive at different conclusions. +" Perlmanetal.(2007) present ground-based Subaru and IRS spectra of the M87 nucleus and find evidence of an excess at wavelengths longer than 25 um, which they attribute to thermal emission from cool dust at a characteristic temperature of some 55 K. This claim is countered by Busonetal. (2009),, who present a higher signal-to-noise IRS spectrum of the nucleus."," \citet{2007ApJ...663..808P} present ground-based Subaru and IRS spectra of the M87 nucleus and find evidence of an excess at wavelengths longer than 25 $\mu$ m, which they attribute to thermal emission from cool dust at a characteristic temperature of some 55 K. This claim is countered by \citet{2009ApJ...705..356B}, who present a higher signal-to-noise IRS spectrum of the nucleus." +" After careful subtraction of a stellar emission template from the mid-infrared spectrum, these authors conclude that the nuclear spectrum can be fully accounted for by optically thin synchrotron emission and that there is little room for dust emission."," After careful subtraction of a stellar emission template from the mid-infrared spectrum, these authors conclude that the nuclear spectrum can be fully accounted for by optically thin synchrotron emission and that there is little room for dust emission." +" On a larger scale, Xilouriset present ISOCAM imaging of M87 and argue that"," On a larger scale, \citet{2004A+A...416...41X} present ISOCAM imaging of M87 and argue that" +Three of the six DAZs with companions have previously measured orbital periods. primary. and companion masses.,"Three of the six DAZs with companions have previously measured orbital periods, primary, and companion masses." + WD 0419-487 is a ML. wwhite dwarf with a 0.095 companion in a 7.3 hr orbit., WD 0419-487 is a $\Msun$ white dwarf with a 0.095 companion in a 7.3 hr orbit. + The spectroscopic survey of DAs that discovered WD 0419-437 noted that there was evidence of a hot spot/accretion disk in the observed spectrum (?).., The spectroscopic survey of DAs that discovered WD 0419-487 noted that there was evidence of a hot spot/accretion disk in the observed spectrum \citep{zuckerman03}. + I would expect that this object may show a higher accretion rate due to some amount of mass {rausfer in addition to that eaused by a wind., I would expect that this object may show a higher accretion rate due to some amount of mass transfer in addition to that caused by a wind. + WD 10264-002 is à 0.68 AL. white chvarl with a 0.23. M. companion in a 14.3 hr orbit., WD 1026+002 is a 0.68 $\Msun$ white dwarf with a 0.23 $\Msun$ companion in a 14.3 hr orbit. + Finally. WD 12134-528 is à 0.63. AL. while cwarl with a 0.36 M. companion in à 16 hr orbit.," Finally, WD 1213+528 is a 0.63 $\Msun$ white dwarf with a 0.36 $\Msun$ companion in a 16 hr orbit." + Table 1. lists the orbital parameters importaul for my ealeulations., Table \ref{tab:closeorb} lists the orbital parameters important for my calculations. + The remaining three svstems have had their radial velocities measured to an accuracy within a few km/s and show either slow RV trends or none al all (2).., The remaining three systems have had their radial velocities measured to an accuracy within a few km/s and show either slow RV trends or none at all \citep{schultz96}. + These WDs are prime candidates for hieh spatial resolution imaging in order to determine what their approximate orbital separations are., These WDs are prime candidates for high spatial resolution imaging in order to determine what their approximate orbital separations are. + Recently. a large snapshot survey (Program 10255. PI D. Που) with the IIubble Space Telescope's Advanced Camera lor Surveys (ACS) was performed on a large sample of previously unresolved white cdwarltred dwarf pairs. including (hie three DAZ-M systems that have no orbital information (forfirstresults.see?)..," Recently, a large snapshot survey (Program 10255, PI D. Hoard) with the Hubble Space Telescope's Advanced Camera for Surveys (ACS) was performed on a large sample of previously unresolved white dwarf+red dwarf pairs, including the three DAZ+M systems that have no orbital information \citep[for first results, see][]{farihi06}." + Of these three. WD 10494-1023 and WD 12102-464 show resolved companions with the ACS data.," Of these three, WD 1049+103 and WD 1210+464 show resolved companions with the ACS data." + The third DAZ. WD 03544463 shows no obvious resolved companion (?)..," The third DAZ, WD 0354+463 shows no obvious resolved companion \citep{farihi06}." + The STSclI pipeline calibrated ACS data lor these three WDs was retrieved aud analvzed bv measuring the centroid positions of the white dwarl and resolved companion., The STScI pipeline calibrated ACS data for these three WDs was retrieved and analyzed by measuring the centroid positions of the white dwarf and resolved companion. + Figure 1 shows the two imaged targets wilh resolved companions., Figure \ref{fig:acs} shows the two imaged targets with resolved companions. + The data was taken with the Ες Resolution Camera (IIRC) on ACS and so the plate seale for the final geometrically corrected images is 070025/pixel., The data was taken with the High Resolution Camera (HRC) on ACS and so the plate scale for the final geometrically corrected images is $\sim$ 025/pixel. + The projected radius was calculated and this was used to estimate H1 orbital separation assuming a circular orbit., The projected radius was calculated and this was used to estimate an orbital separation assuming a circular orbit. + Since no resolved companion was detected for WD 0354-4463. an estimate for the upper limit to the AL dwarl orbital separation must be done.," Since no resolved companion was detected for WD 0354+463, an estimate for the upper limit to the M dwarf orbital separation must be done." + One can estimate the sensitivity of an image to a resolved companion bv implanting artificial companions into the data and measure when (μον are recovered., One can estimate the sensitivity of an image to a resolved companion by implanting artificial companions into the data and measure when they are recovered. + I constructed a reference WRC PSF from the image ol WD 1210--464., I constructed a reference HRC PSF from the image of WD 1210+464. + At separations 270 mas or 2 AU. an object with msi 15.4 could," At separations $>$ 70 mas or 2 AU, an object with $_{F814W}$ =15.4 could" +We have already seen that Nyy=S is a suitable choice. which gives No~10.2.,"We have already seen that $N_{\rm ray} = 8$ is a suitable choice, which gives $\bar{N} \sim 10.2$." + This corresponds to a relative error NOU17 per cont.," This corresponds to a relative error $\bar{N}^{-3/4} +\sim 17$ per cent." + Although we can relax the constraint on the fraction of the source plane covered: somewhat. and still have a well sampled MPLHE with the Ray Bundle method. we do not have this Dexibility with the grid. based Ray Shooting method.," Although we can relax the constraint on the fraction of the source plane covered somewhat, and still have a well sampled MPH with the Ray Bundle method, we do not have this flexibility with the grid based Ray Shooting method." + In addition. as we decrease the source size. the number of pixels required. for the RSAL increases. and a higher density of ravs is necessary.," In addition, as we decrease the source size, the number of pixels required for the RSM increases, and a higher density of rays is necessary." + shows the Alagnilication Probability histograms obtained for a single Schwarzschild lens using the Rav Shooting (thin line) and Ray Bunelle (thick line) methods., shows the Magnification Probability histograms obtained for a single Schwarzschild lens using the Ray Shooting (thin line) and Ray Bundle (thick line) methods. + “Phe parameters for each method are [isted in2., The parameters for each method are listed in. +. Phe vertical axis of this (ancl later histograms) is the normalised number of buncles/sources in each magnification bin. ΑΝ). which is equivalent to the definition of p(yAye in(7)..," The vertical axis of this (and later histograms) is the normalised number of bundles/sources in each magnification bin, $N(\mu)$, which is equivalent to the definition of $p(\mu_1)\Delta\mu$ in." + A ecut-olf was imposed. on image locations for the RBAL at ron=101 Einstein radii, A cut-off was imposed on image locations for the RBM at $x_{\rm cut}=1.01$ Einstein radii. +" The two distributions are qualitatively very similar, when the various caveats described above are considered."," The two distributions are qualitatively very similar, when the various caveats described above are considered." + Imposing a larger value of Hou serves Lo reduce the maximum magnification with the RBAL, Imposing a larger value of $x_{\rm cut}$ serves to reduce the maximum magnification with the RBM. + The poor sampling in the highest so bins for the RSAL is clearly demonstrated., The poor sampling in the highest $\mu$ bins for the RSM is clearly demonstrated. + The dashed. line shows the expected (iT power lw. slope of the MPLL at large jr (see A)). and both distributions have this approximate form (although for the RSAL the statistical significance of the histogram bins with yr>10 is low).," The dashed line shows the expected $\mu^{-2}$ power law slope of the MPH at large $\mu$ (see ), and both distributions have this approximate form (although for the RSM the statistical significance of the histogram bins with $\mu >10$ is low)." + As expected. the RSAL produces magnifications which aro peSL (when fiaupte22ld ds expected) due to numerical cllects.," As expected, the RSM produces magnifications which are $\mu < 1$ (when $\mu_{\rm empty} \geq 1$ is expected) due to numerical effects." + Phe sample mean and variance of the two distributions are (rp=0.98 and ση=0.03 for Ray Shooting. and (ji)=1.02 and o7;=0.08 for the Ray Dundlemethod?.," The sample mean and variance of the two distributions are $\langle\mu\rangle = 0.98$ and $\sigma^2_{\mu} = 0.03$ for Ray Shooting, and $\langle\mu\rangle = 1.02$ and $\sigma^2_{\mu} = 0.08$ for the Ray Bundle." +. The Ray Shooting method provides higher accuracy at high magnifications. but at magnifications yp1. the Rav Bunelle method is more accurate tux from additional images is neglected.," The Ray Shooting method provides higher accuracy at high magnifications, but at magnifications $\mu \sim 1$, the Ray Bundle method is more accurate flux from additional images is neglected." + One aspect of the MPLIL we have not vet discussed has to do with the weighting we apply to each ray bundle., One aspect of the MPH we have not yet discussed has to do with the weighting we apply to each ray bundle. + For the ΛΙ. every source is a pixel with the same area.," For the RSM, every source is a pixel with the same area." + bor the RM. the initial ray bundles (images) have the same area. but the resulting sources must have cdillerent areas (bv the definition of a magnilication).," For the RBM, the initial ray bundles (images) have the same area, but the resulting sources must have different areas (by the definition of a magnification)." + It is sullicient to weight each rav. bundle in the MPII by the area of the resulting source., It is sufficient to weight each ray bundle in the MPH by the area of the resulting source. + shows the RBAI (thin line) with no weighting compared with the correct area weighting (thick line)., shows the RBM (thin line) with no weighting compared with the correct area weighting (thick line). + For an ensemble of No lenses in the lens plane. cach with mass AZ). the total dellection angle generalises to where the (€—£;) are the impact parameters to each lens.," For an ensemble of $N$ lenses in the lens plane, each with mass $M_j$, the total deflection angle generalises to where the $(\bvec{\xi} - \bvec{\xi}_j)$ are the impact parameters to each lens." + Consider the case where lenses are restricted to. [ie within a rectangularregion., Consider the case where lenses are restricted to lie within a rectangular. +. Light ravs passing through one of the corners of the shooting region will necessarily be dellectec inwards bv. the mass. distribution., Light rays passing through one of the corners of the shooting region will necessarily be deflected inwards by the mass distribution. + This. is appropriate for an isolated configuration of lenses. such as in studies of the microlensing ellect of many stars which make up a galaxy (where the contribution to the deflection," This is appropriate for an isolated configuration of lenses, such as in studies of the microlensing effect of many stars which make up a galaxy (where the contribution to the deflection" +In Fig.,In Fig. +" 2, we show the proton and neutron Thomas-Fermi distributions for 7°°Pb as obtained with the method describe above and the DBHF model for the EoS. 'The predicted proton and neutron root-mean-square radii are 5.39 fm and 5.56 fm, respectively."," 2, we show the proton and neutron Thomas-Fermi distributions for $^{208}$ Pb as obtained with the method describe above and the DBHF model for the EoS. The predicted proton and neutron root-mean-square radii are 5.39 fm and 5.56 fm, respectively." +" Typically, predictions of the symmetry energy at saturation density encountered in"," Typically, predictions of the symmetry energy at saturation density encountered in" +The comparison of the CVDs needs some prior comments.,The comparison of the CVDs needs some prior comments. + Generally. our CVD measurements agree with the values given in the above-mentioned works. and with those listed by McElroy (1995). with the exception of NGC 1399. for which McElroy gives a value of 308 kms +.," Generally, our CVD measurements agree with the values given in the above-mentioned works, and with those listed by McElroy (1995), with the exception of NGC 1399, for which McElroy gives a value of 308 km $^{-1}$." + D95 give 420427 kms +. that is still under the 430 km ! derived by Stiavelli et al. (," D95 give $\pm$ 27 km $^{-1}$, that is still under the 430 km $^{-1}$ derived by Stiavelli et al. (" +"1993) using core resolution techniques applied to observations performed under 0.7"" seeing.",1993) using core resolution techniques applied to observations performed under $\arcsec$ seeing. + Our CVDs are systematically lower than those by D95. that were obtained under better seeing conditions (private communication).," Our CVDs are systematically lower than those by D95, that were obtained under better seeing conditions (private communication)." + On the other hand. since the smoothing effect of the seeing is stronger for the steeper VDPs. and since the gradient of the VDP is correlated to the CVD (Busarello et al.," On the other hand, since the smoothing effect of the seeing is stronger for the steeper VDPs, and since the gradient of the VDP is correlated to the CVD (Busarello et al." + 1997). we expect the lower CVDs to be less affected by the smoothing.," 1997), we expect the lower CVDs to be less affected by the smoothing." + The two sets of data are related by: los(od)=L.10(43-.01)«los(oi—0.17(—.1) (see Fig.," The two sets of data are related by: $\log(\sigma_0^{\tiny D95})=1.10(\pm.04)\times \log(\sigma_0^{\tiny +this\ work})-0.17(\pm.1)$ (see Fig." + D., 1). +" We can then use this relation to estimate the ""corrections? to the CVDs for the other objects of our sample.", We can then use this relation to estimate the `corrections' to the CVDs for the other objects of our sample. +" These ""corrected"" CVDs are also listed in Table 2.", These `corrected' CVDs are also listed in Table 2. + Of the 12 ellipticals studied here. five (NGC 1339. NGC 1374. NGC 1379. NGC 1404 and NGC 1419) are rotationally supported systems. having positive rotation parameter + )*0).," Of the 12 ellipticals studied here, five (NGC 1339, NGC 1374, NGC 1379, NGC 1404 and NGC 1419) are rotationally supported systems, having positive rotation parameter $\frac{V}{\sigma}$ $>$ 0)." + The interest in this result is strengthened by the fact that four of these galaxies are EU/EI systems. the only remaining round system in the sample being NGC 1399.," The interest in this result is strengthened by the fact that four of these galaxies are E0/E1 systems, the only remaining round system in the sample being NGC 1399." + NGC 1374. NGC 1379 and NGC 404 are most probably SO galaxies (due to the overall constancy of their RCs: see also Fisher 1997. and the case of NGC 3379 presented in Statler et al.," NGC 1374, NGC 1379 and NGC 1404 are most probably S0 galaxies (due to the overall constancy of their RCs; see also Fisher 1997, and the case of NGC 3379 presented in Statler et al." + 1997). although a more complex structure has to be invoked to explain their kinematical features.," 1997), although a more complex structure has to be invoked to explain their kinematical features." + The RCs of NGC 1374 and NGC 1404. in particular. may reveal the presence of a double-disk structure (cf.," The RCs of NGC 1374 and NGC 1404, in particular, may reveal the presence of a double-disk structure (cf." + Seifert Scorza 1996)., Seifert Scorza 1996). +" The rotation parameter of NGC 1339 ts consistent with that of an isotropic oblate model. but the shape and the asymmetry of both the VDP and the RC indicate a significant departure from a ""simple? isotropic rotator."," The rotation parameter of NGC 1339 is consistent with that of an isotropic oblate model, but the shape and the asymmetry of both the VDP and the RC indicate a significant departure from a `simple' isotropic rotator." + Four objects (NGC 1336. FCC 136. NGC 1373 and NGC 1419) are most probably SBO galaxies: NGC 1336 and NGC 1419. in particular. have kinematical profiles identical to the SBO galaxy IC 456 (Betton! 1989. see also Galletta 1996).," Four objects (NGC 1336, FCC 136, NGC 1373 and NGC 1419) are most probably SB0 galaxies: NGC 1336 and NGC 1419, in particular, have kinematical profiles identical to the SB0 galaxy IC 456 (Bettoni 1989, see also Galletta 1996)." + This conclusion is also strongly supported by the photometric properties. all these galaxies having strong tsophotal twisting (see the photometric profiles in Caon et al.," This conclusion is also strongly supported by the photometric properties, all these galaxies having strong isophotal twisting (see the photometric profiles in Caon et al." + 1994)., 1994). + The three remaining galaxies are definitely “dynamically hot” systems., The three remaining galaxies are definitely `dynamically hot' systems. +" FCC 335 is a faint galaxy (Mp= 16.8 for Hy=75 km ? ‘) with low surface brightness (j/,223.5 B mag/aresec?). low velocity dispersion (σι 43 km 1) and low rotation (V,,,,2 24 km 1j."," FCC 335 is a faint galaxy $_B$ $-$ 16.8 for $_0$ =75 km $^{-1}$ $^{-1}$ ) with low surface brightness $\mu_e$ =23.5 B $^2$ ), low velocity dispersion $\sigma_0$ = 43 km $^{-1}$ ) and low rotation $_{max}$ = 24 km $^{-1}$ )." +" NGC 1399 and NGC 1427 are instead bright galaxies. with high (and roughly constant) velocity dispersions. and relatively slow rotation: this is exactly what one expects from ""true? giant ellipticals."," NGC 1399 and NGC 1427 are instead bright galaxies, with high (and roughly constant) velocity dispersions, and relatively slow rotation: this is exactly what one expects from `true' giant ellipticals." + For both objects the present data show complex kinematics: the already-known kinematically distinct core in NGC 1399 plus a second clear kinematically distinct component, For both objects the present data show complex kinematics: the already-known kinematically distinct core in NGC 1399 plus a second clear kinematically distinct component +Galaxies are drawn from the Sloan Digital Sky Survey Data Release 7 ο.DIU.Abazajianetal..2009..,"Galaxies are drawn from the Sloan Digital Sky Survey Data Release 7 \citealp[SDSS DR7,][]{dr7}." + In particular. we make use of the publicly available New York Value Added: Catalog (NYU-VAC) released by Blantoneal. (2005).," In particular, we make use of the publicly available New York University-Value Added Catalog (NYU-VAC) released by \citet{blanton05}." +. We select. galaxies in the redshift range 0.04.< and masses log Ad.οLOA..., We select galaxies in the redshift range $0.0410.4 M_\odot$. +" Phe magnitude limi in the SDSS spectroscopy places the lower mass limit for red galaxies at z0.1 approximately at logM,;=10.6M.. although we choose to work with galaxies slightly below tha limit."," The magnitude limit in the SDSS spectroscopy places the lower mass limit for red galaxies at $z\sim 0.1$ approximately at $M_*=10.6 M_\odot$, although we choose to work with galaxies slightly below that limit." +" Nonetheless. we warn the reader that our sample is incomplete by ~20% below log Al,=10.6M at the highes redshift probed here."," Nonetheless, we warn the reader that our sample is incomplete by $\sim$ below log $M_*=10.6 M_\odot$ at the highest redshift probed here." + We make use of NYU-VAC. k-correcte photometry. spectroscopic redshifts ancl light-profile fitting parameters in the r band.," We make use of NYU-VAC k-corrected photometry, spectroscopic redshifts and light-profile fitting parameters in the r band." +" Stellar masses. metallicities anc r-band. light-weighte ages [or SDSS DIU galaxies have been estimated. in the same way as for previous releases and as described in G05, o which we refer the reader for a full description."," Stellar masses, metallicities and r-band light-weighted ages for SDSS DR7 galaxies have been estimated in the same way as for previous releases and as described in G05, to which we refer the reader for a full description." + Drielly. estimates of the stellar population parameters are obtained » comparing the observed. stellar absorption features (corrected. for emission. lines) with those predicted. by a comprehensive library of model spectra based. on Bruzual&Charlot(2003) SSPs convolved with Monte Carlo star ormation histories.," Briefly, estimates of the stellar population parameters are obtained by comparing the observed stellar absorption features (corrected for emission lines) with those predicted by a comprehensive library of model spectra based on \citet{bc03} SSPs convolved with Monte Carlo star formation histories." + A comparison between the new SDSS DIU parameters and those of SDSS DRA from (105 and C06 provides no systematic olfset and a typical dispersion at the level of —0.1 dex in light-weighted age. stellar mass ancl metallicity for galaxies with red. stellar populations. well below the typical error. budget in. those measurements.," A comparison between the new SDSS DR7 parameters and those of SDSS DR4 from G05 and G06 provides no systematic offset and a typical dispersion at the level of $\sim 0.1$ dex in light-weighted age, stellar mass and metallicity for galaxies with red, stellar populations, well below the typical error budget in those measurements." + In this work we correct the estimated galaxy ages to z=0 bv adding the lookback time at the redshift of he galaxy under the assumption of passive evolution (which is very reasonable for our sample of quiescent galaxies), In this work we correct the estimated galaxy ages to z=0 by adding the lookback time at the redshift of the galaxy under the assumption of passive evolution (which is very reasonable for our sample of quiescent galaxies). + In addition to the aforementioned. stellar. population oxwanmeters we focus on the a enhancement., In addition to the aforementioned stellar population parameters we focus on the $\alpha-$ enhancement. + In spite of he name. the effect. is more a lack of iron rather than an excess of @ elements.," In spite of the name, the effect is more a lack of iron rather than an excess of $\alpha-$ elements." + This lack of iron. is produced when the SE timescale of à galaxy. is short., This lack of iron is produced when the SF timescale of a galaxy is short. + Core-collapse supernovae enrich the medium with a-elements in scales of a lew tenths of Myrs. while the Fe-enrichmoent is due to tvpe la supernovae explosions.," Core-collapse supernovae enrich the medium with $\alpha$ -elements in scales of a few tenths of Myrs, while the Fe-enrichment is due to type Ia supernovae explosions." + Hf a significant fraction of the stars ave formed in a period shorter than the ~IE Civr needed by ἵνρο la supernovae progenitors to evolve. the stars would show a chemical composition with higher α ο abundance ratios than those in stellar populations formed. with longer timescales.," If a significant fraction of the stars are formed in a period shorter than the $\sim 1$ Gyr needed by type Ia supernovae progenitors to evolve, the stars would show a chemical composition with higher $\alpha$ /Fe abundance ratios than those in stellar populations formed with longer timescales." + As a tracer of the a-cnhancement in the stellar populations of our sample we use the semi-enipirical adpha/Ee] indicator adopted by GOG. namely AC(Mgb/ej) which is the cdillerence. between the observed Mgb/£Fe absorption and that of the scaled-solar BCO3 model that best fits α EFe]-independent features.," As a tracer of the $\alpha$ -enhancement in the stellar populations of our sample we use the semi-empirical [alpha/Fe] indicator adopted by G06, namely $\Delta$ $\langle$ $\rangle$ ) which is the difference between the observed $\langle$ $\rangle$ absorption and that of the scaled-solar BC03 model that best fits $\alpha$ /Fe]-independent features." + G06 have tested. through comparison with Thomas et al. (," G06 have tested, through comparison with Thomas et al. (" +2003) models with variable abundance ratios. that ;NCMgb/iEe3)) correlates linearly. with the abundance ratio α Κο) independently of age and metallicity (except. for metallicities below 30 percent solar. which is lower than the range covered by our sample).,"2003) models with variable abundance ratios, that $\Delta$ $\langle$ $\rangle$ )) correlates linearly with the abundance ratio $\alpha$ /Fe] independently of age and metallicity (except for metallicities below 30 percent solar, which is lower than the range covered by our sample)." + In particular. we confirm such proportionality over the metallicity and age range spanned by our sample (0.5«Z/Z.2 and age older than 3 Giver) with the differential models presented in Walcheretal.(2009). based on the theoretical Coelhoetal.(2007) models calibrated onto either BCOS or Vazdekisetal.(2010).," In particular, we confirm such proportionality over the metallicity and age range spanned by our sample $0.5 AU) separations."," 1999), the IMF of low-mass companions is not well understood, especially at “wide” $>$ AU) separations." + Radial velocity searches around solu-tvpe nail sequence stars (e.g. Mavor Queloz 1995: Marev Butler 1996) have produced CW COwired brown clwarts at separations <33 AAU.," Radial velocity searches around solar-type main sequence stars (e.g., Mayor Queloz 1995; Marcy Butler 1996) have produced few confirmed brown dwarfs at separations $<$ AU." +" Fewer tha10.5% of their sample wave brown dwurf conmpanious at those separatioIs,", Fewer than of their sample have brown dwarf companions at those separations. + A coronagraphlic search for compatnicuns in the range LOOAAT (Oppenheimer et al., A coronagraphic search for companions in the range AU (Oppenheimer et al. +" 200]) produced. only one own dwarf. CJ 229B (Naka,Jima et al."," 2001) produced only one brown dwarf, GJ 229B (Nakajima et al." + 1995). well below he multiplicity obscvved for all stars (Reid Cagis 1997).," 1995), well below the multiplicity observed for all stars (Reid Gizis 1997)." + Other tvpes of surveys. such as high spatial resolution space-based observations (Lowrance et al.," Other types of surveys, such as high spatial resolution space-based observations (Lowrance et al." + 1999: Lowrance et al., 1999; Lowrance et al. + 2000) and erotud-sed adaptive optics (Els et al., 2000) and ground-based adaptive optics (Els et al. + 2001). have also resulted in discoveries of low-mass stellar ando sub-stellar companions.," 2001), have also resulted in discoveries of low-mass stellar and sub-stellar companions." +" Towever. the frequency of stellar and suh-stellar conipauious at close separations retains distinctly different. resulting iu the idea that there is a ""brown dwarf desert”."," However, the frequency of stellar and sub-stellar companions at close separations remains distinctly different, resulting in the idea that there is a “brown dwarf desert”." + To date there has 30011 ouly OC SVScluatic search for brown dwarf conipaniois af wide separations and with a volume-limited sanipe (Simons ct al., To date there has been only one systematic search for brown dwarf companions at wide separations and with a volume-limited sample (Simons et al. + 1996: hereafter. STIs).," 1996; hereafter, SHK)." + This was niaiilv à color-based search around M. dwarts within ppc of: the Sun aud did not turi up anv new brown dwarfs. athough. given the surprisingly blue colors of C.J 229B. coo brown dwarfs with intermediate J-K colors may have been overlooked 1i the survey.," This was mainly a color-based search around M dwarfs within pc of the Sun and did not turn up any new brown dwarfs, although, given the surprisingly blue colors of GJ 229B, cool brown dwarfs with intermediate $J$ $K$ colors may have been overlooked in the survey." + Proper motion searches for companions have been used for many wears to ideiifv low-mass objects (e.g. vau Diesbroeck 1961) iux oftor a less biased wav of finding low-lass conipauions than color-based surveys.," Proper motion searches for companions have been used for many years to identify low-mass objects (e.g., van Biesbroeck 1961) and offer a less biased way of finding low-mass companions than color-based surveys." + ThereOre. Wwe have couducted the planned second epoch survey of the SII. sample. im oxer to identify low-mass comMOMS to M dviufs at wide separations out to over AAT," Therefore, we have conducted the planned second epoch survey of the SHK sample, in order to identify low-mass companions to M dwarfs at wide separations out to over AU." +", The choice of M «buf prinaerles is sienificaut: Reid Cagis (1997) ane Reid et al. (", The choice of M dwarf primaries is significant: Reid Gizis (1997) and Reid et al. ( +1999) show that the distribution of dass rajos for a suuple of AL dwarts has a peak at q 095. where q is the ratio of the secondary mass to f1C ordnuarvonanss.,"1999) show that the distribution of mass ratios for a sample of M dwarfs has a peak at $q=$ 0.95, where $q$ is the ratio of the secondary mass to the primary mass." + They couchde that their sample shows a distinct bias towards approximately οςpualanass αν{ος alc that the mass fiction for stellar colpanious ids different from the IME of fold stars., They conclude that their sample shows a distinct bias towards approximately equal-mass systems and that the mass function for stellar companions is different from the IMF of field stars. + Tf these couchisions extend to brown dwarf masses. A primaries may harbor nore sub-stellar coiipauions than other stellar types.," If these conclusions extend to brown dwarf masses, M primaries may harbor more sub-stellar companions than other stellar types." + On he other haud. Reipurth Clarke," On the other hand, Reipurth Clarke" +al loci deep within the cloud. which may approach in dense cores.,"at loci deep within the cloud, which may approach in dense cores." + We estimate the timescale for CO depletion in a core of density my~10!em? to be ~0.6 Myr. a value consistent with independent estimates of the lifetime of a typical core.," We estimate the timescale for CO depletion in a core of density $n_{\rm H}\sim 10^4~{\rm cm}^{-3}$ to be $\sim 0.6$ Myr, a value consistent with independent estimates of the lifetime of a typical core." + Dispersal of cores during star formation mav help to maintain observable levels of gaseous CO over cloud lifetimes., Dispersal of cores during star formation may help to maintain observable levels of gaseous CO over cloud lifetimes. + Our results will provide useful constraints on astrochemical models for molecular clouds (e.g.. Chang 22007; Cuppen Herbst 2007). and on estimates of macroscopic cloud properties such as mass. structure. and the distribution of molecular material (e.g.. Padoan 22006: Pineda 22010).," Our results will provide useful constraints on astrochemical models for molecular clouds (e.g., Chang 2007; Cuppen Herbst 2007), and on estimates of macroscopic cloud properties such as mass, structure, and the distribution of molecular material (e.g., Padoan 2006; Pineda 2010)." + The constraints placed on CO depletion bx our analysis probably represent the limits of what can be accomplished by combining data from the disparate techniques of gas-phase enission spectroscopy. and solid-state absorption spectroscopy., The constraints placed on CO depletion by our analysis probably represent the limits of what can be accomplished by combining data from the disparate techniques of gas-phase emission spectroscopy and solid-state absorption spectroscopy. + Further refinement awails a Largeted survey of inlrared. vibration-rotation lines in absorption toward background field stars known to display ice features. requiring sspectra of high signal-to-noise with resolving powers Z101.," Further refinement awaits a targeted survey of infrared vibration-rotation lines in absorption toward background field stars known to display ice features, requiring spectra of high signal-to-noise with resolving powers $\ga 10^4$." + We also advocate further study of intermecdiate-1nnass YSOs by (his technique. to investigate the degree to which Chev may contribute to the global recvcling of CO between solid and gaseous pliases in molecular clouds. and to better evaluate other astrochenical pathways for CO. such as CIIOIIL formation. in protostellar envelopes.," We also advocate further study of intermediate-mass YSOs by this technique, to investigate the degree to which they may contribute to the global recycling of CO between solid and gaseous phases in molecular clouds, and to better evaluate other astrochemical pathways for CO, such as $_3$ OH formation, in protostellar envelopes." + This research has made use of data from the Two Micron All Sky Survey. which is a joint project of the University of Massachusetts and the Infrared Processing and Ànalvsis Center. [unded by the National Aeronautics and Space Acministration (NASA) and the National Science Foundation.," This research has made use of data from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center, funded by the National Aeronautics and Space Administration (NASA) and the National Science Foundation." + aacknowledges financial support from the NASA Exobiology and Evolutionary Biology program (grant NNXOTAIN38C) and the NASA Astrobiology Institute (grant. NNAO9DASOA)., acknowledges financial support from the NASA Exobiology and Evolutionary Biology program (grant NNX07AK38G) and the NASA Astrobiology Institute (grant NNA09DA80A). + This work was supported in part by the Jet Propulsion Laboratory. California Institute of Technology. under a contract with NASA.," This work was supported in part by the Jet Propulsion Laboratory, California Institute of Technology, under a contract with NASA." + We are grateful lo an anonvimous referee for helpful comments., We are grateful to an anonymous referee for helpful comments. +the rratio to correct the data for reddening and obtained 1 1)) = 0.54.,the ratio to correct the data for reddening and obtained $C$ ) = 0.54. +" With the help of tasks in the ""nebular packagee. we compute Zi = 15700 Ix. using the 1u]]. AbOUT /A4363 ratio."," With the help of tasks in the `nebular package', we compute $T_{\rm e}$ = 15700 K, using the ] $\lambda$ $\lambda$ 4363 ratio." +" Assuming an electron density of ος = 100 "" we compute the tonic abundances for and and derive a total oxvecn abundance of 12|log(O/I11) = 7.95 (—0.2 Z.)). typical for BCD egalaxics."," Assuming an electron density of $n_{\rm e}$ = 100 $^{-3}$ we compute the ionic abundances for $^+$ $^+$ and $^{++}$ $^+$ and derive a total oxygen abundance of 12+log(O/H) = 7.95 $\sim$ ), typical for BCD galaxies." + The large metallicity dillerence between NGC 1512 and NGC 1510 indicates that both galaxies have experienced a very diferent chemical evolution. and that NGC 1510 has been in a quiescent state for a long time while NGC 1512 was forming stars continuously.," The large metallicity difference between NGC 1512 and NGC 1510 indicates that both galaxies have experienced a very different chemical evolution, and that NGC 1510 has been in a quiescent state for a long time while NGC 1512 was forming stars continuously." + The N/O ratio found in NGC 1510. log(N/O) = is rather high for a galaxy with its oxvgen abuncance.," The N/O ratio found in NGC 1510, log(N/O) $\approx$ --1.2, is rather high for a galaxy with its oxygen abundance." + For comparison. Izotoy Thuan. (2004) typically obtain log(N/O) z 1.5 for low metallicity BCD galaxies.," For comparison, Izotov Thuan (2004) typically obtain log(N/O) $\approx$ –1.5 for low metallicity BCD galaxies." + Similar results are also found for other DCD galaxies with a significant population of WI stars (Brinchmann et al., Similar results are also found for other BCD galaxies with a significant population of WR stars (Brinchmann et al. + 2008)., 2008). + Lt is thought that the nitrogen enrichment is à consequence ofa very recent chemical pollution event probably connected with the onset of WI winds (Lóppez-Sánnchez et al., It is thought that the nitrogen enrichment is a consequence of a very recent chemical pollution event probably connected with the onset of WR winds (Lóppez-Sánnchez et al. + 2007)., 2007). + ]t max. also be related to the interaction between galaxies (Pustilnik et al., It may also be related to the interaction between galaxies (Pustilnik et al. + 2004. Lóppez-Sánnchez et al.," 2004, Lóppez-Sánnchez et al." + 2009). but deeper optical spectroscopic data with a higher spectral resolution are needed to confirm this issue.," 2009), but deeper optical spectroscopic data with a higher spectral resolution are needed to confirm this issue." + Star formation depends on the gravitational instability of ealaxy disks. both locally and globally.," Star formation depends on the gravitational instability of galaxy disks, both locally and globally." + Minor mergers and tidal interactions allect the eas distribution and. cdvnanmies of galaxies. leading to the formation of bars. gas inflow as well as the ejection of gas. ancl as a consequence — locally enhanced star formation.," Minor mergers and tidal interactions affect the gas distribution and dynamics of galaxies, leading to the formation of bars, gas inflow as well as the ejection of gas, and – as a consequence – locally enhanced star formation." + Together. these phenomena are kev ingredients to the understanding of galaxy evolution.," Together, these phenomena are key ingredients to the understanding of galaxy evolution." + The often extended cenvelopes of spiral galaxies are particularly useful as sensitive tracers of tidal interactions and gas accretion., The often extended envelopes of spiral galaxies are particularly useful as sensitive tracers of tidal interactions and gas accretion. + The gas distribution and dvnanmies are easily inlluenced by the environment. resulting in asvmimetrics. linc broadening and/or splitting etc.," The gas distribution and dynamics are easily influenced by the environment, resulting in asymmetries, line broadening and/or splitting etc." + The development of a strong two-armed spiral pattern ancl star-forming regions in disk galaxies (here NGC 1512) which accrete low-mass dwarf companions (here NGC 1510) has been explored by Mihos Lernquist (1994). using numerical simulations., The development of a strong two-armed spiral pattern and star-forming regions in disk galaxies (here NGC 1512) which accrete low-mass dwarf companions (here NGC 1510) has been explored by Mihos Hernquist (1994) using numerical simulations. + Fheir models. which use a mass ratio of 10:1 for the disk galaxy ancl its companion. resemble the galaxy pair. NGC 1512/1510 after ~40 time units (1.6. 5.2107 vers).," Their models, which use a mass ratio of 10:1 for the disk galaxy and its companion, resemble the galaxy pair NGC 1512/1510 after $\sim$ 40 time units (i.e. $5.2 \times 10^8$ years)." + At that stage. the model. disk. galaxy has developed a pronounced. slightly asvmamoetrie two-armed spiral pattern with significant star-formation along the armis and the nuclear region.," At that stage, the model disk galaxy has developed a pronounced, slightly asymmetric two-armed spiral pattern with significant star-formation along the arms and the nuclear region." + Minor mergers are common., Minor mergers are common. + The Milky Way and. the Andromeda galaxy are prominent examples: both have many satellites and show evidence for continuous accretion of small companions., The Milky Way and the Andromeda galaxy are prominent examples; both have many satellites and show evidence for continuous accretion of small companions. + Ehe multitude of stellar streams detected. in our Galaxy as well as some other galaxies (e.g... NGC 5907. Martinez-Delgado et al.," The multitude of stellar streams detected in our Galaxy as well as some other galaxies (e.g., NGC 5907, Martinez-Delgado et al." + 2008) are hinting at a rich accretion history., 2008) are hinting at a rich accretion history. + Minor mergers contribute significantly to galaxy assembly. aceretion. and evolution.," Minor mergers contribute significantly to galaxy assembly, accretion, and evolution." + \\Ὁ analysedIvsed the distributiondbut andLL kinematics| of[the the geas as well as the star formation activity in the galaxy. pair NGC 1512/1510 and its surroundings., We analysed the distribution and kinematics of the gas as well as the star formation activity in the galaxy pair NGC 1512/1510 and its surroundings. + For the barred. double-ring galaxy NGC 1512 we find a very large deisk. about four times the optical diameter. with two »onounced. spiral arms. possibly tically induced. by the interaction with the neighbouring blue compact ναί galaxy NGC 1510.," For the barred, double-ring galaxy NGC 1512 we find a very large disk, about four times the optical diameter, with two pronounced spiral arms, possibly tidally induced by the interaction with the neighbouring blue compact dwarf galaxy NGC 1510." + It is possible that the interaction also rigecred the formation of the bar in NGC 1512. (unless he bar already existed. maybe from a previous accretion or interaction event) which would then cause gas to fall towards he nuclear regions. feeding the star formation. as well as induces torques in the outer spiral armis.," It is possible that the interaction also triggered the formation of the bar in NGC 1512 (unless the bar already existed, maybe from a previous accretion or interaction event) which would then cause gas to fall towards the nuclear regions, feeding the star formation, as well as induces torques in the outer spiral arms." +" We detect two with mumasses of <10"" aand clear signs of star formation in the outer-most regions of the system.", We detect two with masses of $\la10^7$ and clear signs of star formation in the outer-most regions of the system. + The most distant PDC. NI1512-west. is rather compact and lies at a distance of 80 kpe from the centre of NGC 1512. potentially at the tip of an polated eastern aarm of NGC 1512.," The most distant TDG, N1512-west, is rather compact and lies at a distance of $\sim$ 80 kpc from the centre of NGC 1512, potentially at the tip of an polated eastern arm of NGC 1512." + Phe second PDC. 1512-south. is forming within an extended ccloud. and is located: slightly. closer (64 kpe). within the extrapolated eastern aan.," The second TDG, 1512-south, is forming within an extended cloud, and is located slightly closer (64 kpc), within the extrapolated eastern arm." + We regard these two PDCs as typical with respect to their mumass. star forming activity and detachment [rom the interacting svstem.," We regard these two TDGs as typical with respect to their mass, star forming activity and detachment from the interacting system." + While TDCs are often found in major mergers. we find that they can form in mildly interacting system such as NGC 1512/1510.," While TDGs are often found in major mergers, we find that they can form in mildly interacting system such as NGC 1512/1510." + In this case. the interaction is clleetly an aecretion. of a blue. compact dwarf. galaxy (NGC 1510) by the large spiral galaxy NGC 1512.," In this case, the interaction is effectly an accretion of a blue compact dwarf galaxy (NGC 1510) by the large spiral galaxy NGC 1512." + NGC 1512 hosts an extended CV. disk with z200 of clusters with recent star formation activity., NGC 1512 hosts an extended $UV$ disk with $\ga$ 200 of clusters with recent star formation activity. + The comparison ob our nunap with the GALEN images clearly shows that these clumps are located within the maxima of neutral gas density., The comparison of our map with the GALEX images clearly shows that these clumps are located within the maxima of neutral gas density. + We have derived the ages and star formation rates of the UV ιο clusters., We have derived the ages and star formation rates of the $UV$ -rich clusters. + We find that generally only the voungest CV. clusters are associated with high ccolumn densities. while in older CV. clusters only cilfuse σσας is detected.," We find that generally only the youngest $UV$ clusters are associated with high column densities, while in older $UV$ clusters only diffuse gas is detected." + This might suggest that as the hydrogen gas depleted. star formation stopped in the latter regions.," This might suggest that as the hydrogen gas depleted, star formation stopped in the latter regions." +(Herbie1960).. (Caiuinetal.2001).. Porter&Rivinius2003)j. (Beaulicu (Ct,"\citep{her60}, \citep{gri94,wat98,oud01}. \citealt{por03}) \citep{bea01,dew05}." +onzalez—(1997:2005) (Moevuetetal.199LChiappini20," \citep{gon97,fis05} \citep{mey94,chi06}. \citep{wis07}." +06).. deWitetal.2005)) (Deaulieu," \citealt{dew03} \citealt{dew05}) \citep{bea01}. \citep{rob08,urq10}," +etal.2001).. (Robitaille," \citep{lam99,bea01,dew02,dew03} \citep{whi08,cla10}." +well at low redshift.,well at low redshift. + While these models fit the observed relations well. the total mass density of SMDIIS is still well below the observed value.," While these models fit the observed relations well, the total mass density of SMBHs is still well below the observed value." +" Surprisingly. we see nuuginallv steeper slopes in the ""os"" models than iu the ""dt models."," Surprisingly, we see marginally steeper slopes in the “log” models than in the “dt” models." +" Even though ""dt models have higher merger rates aud produce more nassive black holes. they also geuerate a large population of imoderate-uass black holes hosted in small halos."," Even though “dt” models have higher merger rates and produce more massive black holes, they also generate a large population of moderate-mass black holes hosted in small halos." + This skews the fits to flatter slopes despite the higher uaxinmni black hole mass., This skews the fits to flatter slopes despite the higher maximum black hole mass. + However. these differciuces are rot statistically siguificaut. except for the constant imitial nass combined with instant mcreie scenarios. where the ow scatter produces sinall uncertainties.," However, these differences are not statistically significant, except for the constant initial mass combined with instant merging scenarios, where the low scatter produces small uncertainties." +" Generally. the difference between a particular model aud the Af,oo relation is matched by the difference vetween that same imocel aud the Ay,Adar relation."," Generally, the difference between a particular model and the $\msig$ relation is matched by the difference between that same model and the $\mhalo$ relation." + However. the models are more indistinguishable from he AdjMor relation. mostly due to the lareer uncertainties iu the observed relation.," However, the models are more indistinguishable from the $\mhalo$ relation, mostly due to the larger uncertainties in the observed relation." + We have developed a new fast. parallel halo f&uder for inclusion iu cosmological simulations with the simulation code ΕΤΑΡΗ.," We have developed a new fast, parallel halo finder for inclusion in cosmological simulations with the simulation code FLASH." + Using SO halo finding techuiques. we are able to produce halo catalogs iu good agreement with traditional post-processing halo finders.," Using SO halo finding techniques, we are able to produce halo catalogs in good agreement with traditional post-processing halo finders." + Since our halo finder is designed to be fast. we are able to perform halo finding operations at every time step iu the simulation. allowing us to perform a detailed analysis of SMDII suberid models.," Since our halo finder is designed to be fast, we are able to perform halo finding operations at every time step in the simulation, allowing us to perform a detailed analysis of SMBH subgrid models." +" While merging alone caunot senerate chough total nass in SMDIIS to match the observed cosmdüc mass density or generate hieh enough maxinuun black hole mass in the largest halos to match the observed Af,60 and AM,Mor relations. it can plav a large role in developing the slope of the relations. expecially at intermediate mass ranges."," While merging alone cannot generate enough total mass in SMBHs to match the observed cosmic mass density or generate high enough maximum black hole mass in the largest halos to match the observed $\msig$ and $\mhalo$ relations, it can play a large role in developing the slope of the relations, especially at intermediate mass ranges." + Thus. mereing should not be totally discounted in considering the processes that provide the correlations between black hole mass and bulee. galaxy. aud halo properties.," Thus, merging should not be totally discounted in considering the processes that provide the correlations between black hole mass and bulge, galaxy, and halo properties." + However. since none of our considered models cau account for the observed cosmic mass density. there is still a significaut role for aceretion and feedback processes iu the evolution of SMBUs.," However, since none of our considered models can account for the observed cosmic mass density, there is still a significant role for accretion and feedback processes in the evolution of SMBHs." + Also. since the choice of models can greatly affect the cosmüc SAIBIT mass deusitv. accretion aud feedback: models. iust be chosen carefully to match observations.," Also, since the choice of models can greatly affect the cosmic SMBH mass density, accretion and feedback models must be chosen carefully to match observations." + The choice of suberid models can dramatically iupact the mereius rate of black holes., The choice of subgrid models can dramatically impact the merging rate of black holes. + Since the merger rate has a laree influence ou the performance of wpcomine eravitational wave detectors. halo finding operations m sinulations should be done as frequently as possible iu order to accurately capture this rate.," Since the merger rate has a large influence on the performance of upcoming gravitational wave detectors, halo finding operations in simulations should be done as frequently as possible in order to accurately capture this rate." + We do not believe hat the inclusion of gas will sienificautly affect these relative differences. since they are largely driveu by the ability to find more black holes at early times.," We do not believe that the inclusion of gas will significantly affect these relative differences, since they are largely driven by the ability to find more black holes at early times." + However. he choice of mereie test does not ereatlv affect the oxedieted rate. except at low redshift.," However, the choice of merging test does not greatly affect the predicted rate, except at low redshift." + While we have bracketed the possible suberid models with “optimistic” aud “pessimistic” scenarios. models vest matchiug insights from theory aud observations are usually in between those extremes.," While we have bracketed the possible subgrid models with “optimistic” and “pessimistic” scenarios, models best matching insights from theory and observations are usually in between those extremes." + While seeding black oles with a uniforiu mitial mass for black holes may well nodel ligh+vedshitt behavior. it is not clear that this is a useful strategy for re-sccding low-redshitt alos.," While seeding black holes with a uniform initial mass for black holes may well model high-redshift behavior, it is not clear that this is a useful strategy for re-seeding low-redshift halos." + Most re-seeding is certainly an artifact of the halo fiuder aud the ack of au iu-code merecr tree to determine when halos wave truly merged., Most re-seeding is certainly an artifact of the halo finder and the lack of an in-code merger tree to determine when halos have truly merged. + However. there are some plausible scenarios where re-seediug may be needed: for example. when three-body or gas interactions strip an SAIBIT frou a central galaxy.," However, there are some plausible scenarios where re-seeding may be needed: for example, when three-body or gas interactions strip an SMBH from a central galaxy." + Also. when this approach is coupled with infrequent halo finding operations. it may deposit too ittle mass in the seed SMDIIs.," Also, when this approach is coupled with infrequent halo finding operations, it may deposit too little mass in the seed SMBHs." + While instant mcereiue is too optimistic. our current lack of uuderstaudius of SMIDIT mergers nuplies that we cannot cutirely specify lis portion of the suberid model. and we must rely on a bracketing procedure.," While instant merging is too optimistic, our current lack of understanding of SMBH mergers implies that we cannot entirely specify this portion of the subgrid model, and we must rely on a bracketing procedure." + Future examination of these suberid imeocels iuust o0 done iu a cosmological simulation involving eas evolution. accretion onto the black holes. aud feedback roni active ealactic nuclei.," Future examination of these subgrid models must be done in a cosmological simulation involving gas evolution, accretion onto the black holes, and feedback from active galactic nuclei." + However. since models of hese processes carry with them their own assumptions aud adjustable parameters. care must be taken to fully separate the effects of black hole seeding aud mereiue.," However, since models of these processes carry with them their own assumptions and adjustable parameters, care must be taken to fully separate the effects of black hole seeding and merging." + It mav be possible that duc to the sclfvceulating nature of feedback that the differcuces ;nong these models may disappear: however. the differences iu merger rate and peak black hole mass suggest siguificant variances niu remain.," It may be possible that due to the self-regulating nature of feedback that the differences among these models may disappear; however, the differences in merger rate and peak black hole mass suggest significant variances may remain." + Ouly once all aspects of these suberid models are analyzed. understood. aud compared to our observational uiderstaudie cau we confidently combine them iuto au integrated model.," Only once all aspects of these subgrid models are analyzed, understood, and compared to our observational understanding can we confidently combine them into an integrated model." + The authors acknowledge support under a Presideutial Early Career Award from the U.S. Departineut of Encreyv. Lawrence Livermore National Laboratory (contract B532720).," The authors acknowledge support under a Presidential Early Career Award from the U.S. Department of Energy, Lawrence Livermore National Laboratory (contract B532720)." + Additional support was provided by a DOE Computational Scieuce Craduate Fellowship (DE-FO02-97ER25308) and the National Ceuter for Supercolputing Applications., Additional support was provided by a DOE Computational Science Graduate Fellowship (DE-FG02-97ER25308) and the National Center for Supercomputing Applications. + The software used in this work was in part developed by the DOE-supported ASC/Alliauce Ceuter for Astrophysical Thermonuclear Flashes at the University. of Chicago., The software used in this work was in part developed by the DOE-supported ASC/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago. + This research used resources of the National Center for Computational Sciences at Oak. Ridge National Laboratory. which is supported by the Office of Science of the US Departineut of Encrev under contract no.," This research used resources of the National Center for Computational Sciences at Oak Ridge National Laboratory, which is supported by the Office of Science of the US Department of Energy under contract no." + DE-AC'05-000R22725., DE-AC05-00OR22725. + The authors thank Brian O'Shea. Tiziana Di Matteo. aud many others for useful couversatious at the tenth Cxeat Lakes Cosmology. Workshop.," The authors thank Brian O'Shea, Tiziana Di Matteo, and many others for useful conversations at the tenth Great Lakes Cosmology Workshop." +between the core emission and the region of possible infall is further evidence of a collapsine. star forming. core (Naravan,"between the core emission and the region of possible infall is further evidence of a collapsing, star forming, core \citep{nmwb}." +anrelal.2002).. Figure 14. shows the integrated intensity map ol the CO 0) line wings. and we see the outflow in the same direction as detected by Myersοἱal.(1983).," Figure \ref{sfoxviii_outflow} shows the integrated intensity map of the CO ) line wings, and we see the outflow in the same direction as detected by \citet{mhsg}." +. SFO 20 is a tvpe C bright rimmed cloud., SFO 20 is a type C bright rimmed cloud. + The long tail is visible in the optical image (Figure 2)) pointing northwards. away trom the III region. IC434. which is approximately 400 parsecs [rom us (SEO).," The long tail is visible in the optical image (Figure \ref{sfo_dss}) ) pointing northwards, away from the HII region IC434, which is approximately 400 parsecs from us \citetext{SFO}." +. The infrared spectral οποιον distribution of this source is consistent with an embedded YSO (SEO)... Sugi, The infrared spectral energy distribution of this source is consistent with an embedded YSO \citetext{SFO}. +tanietal.(1989). discovered. and mapped an outllow around the embedded IRAS source in SFO 20 using the NRO 45m telescope in ihe CO 0) transition., \citet{sfmo} discovered and mapped an outflow around the embedded IRAS source in SFO 20 using the NRO 45m telescope in the CO ) transition. + The CO 0) and CO 1) line profiles shown in Figure 15aa show a blue asymmetry and wings on the transition., The CO ) and CO ) line profiles shown in Figure \ref{sfoxx_xxv_profiles}a a show a blue asymmetry and wings on the transition. + The ]lime profiles are all fairly. eaussian. and lairiy well centered on the 0) Vin.," The line profiles are all fairly gaussian, and fairly well centered on the ) $V_{\rm LSR}$." + There is no clear evidence of infall from the (transitions., There is no clear evidence of infall from the transitions. + The CO 0) integrated intensity map (Figure 16)) shows the tvpe C morphology. with a strong core and evidence of a tail trailing off to the north.," The CO ) integrated intensity map (Figure \ref{sfoxx_itty}) ) shows the type C morphology, with a strong core and evidence of a tail trailing off to the north." + The CO I) enmission also (traces the core. though the emission peak does not match the 2) core position exactly.," The CO ) emission also traces the core, though the emission peak does not match the ) core position exactly." + The 0) emission is also peaked on the core and relatively absent in the tail., The ) emission is also peaked on the core and relatively absent in the tail. + The 0) emission is mainly centered on the core Wilh some emission tracing (he tail., The ) emission is mainly centered on the core with some emission tracing the tail. + While many of the previous sources had substantil emission outside of the core region. this one does not.," While many of the previous sources had substantial emission outside of the core region, this one does not." + Under the assumption that the tvpe C morphology is (he most temporally evolved morphology. the ionization front may have long since passed this region and it is surrounded by ionized gas.," Under the assumption that the type C morphology is the most temporally evolved morphology, the ionization front may have long since passed this region and it is surrounded by ionized gas." + As the core gas collapses the tail can be eaten away as it becomes less and less shielded., As the core gas collapses the tail can be eaten away as it becomes less and less shielded. + The 2) traces onlv the core material., The ) traces only the core material. + Analvsis of the 0) centroid velocity. shown in Figure 17. indicates a small. but statistically significant. gradient across the core.," Analysis of the ) centroid velocity, shown in Figure \ref{sfoxx_centroid} indicates a small, but statistically significant gradient across the core." + This gradient. may indicate the cloud/s rotation., This gradient may indicate the cloud's rotation. + As an outflow has already been detected in this source by we chose to look at the CO 1) wing integrated intensitv in order (o (rv to verily the existing observation., As an outflow has already been detected in this source by \citet{sfmo} we chose to look at the CO ) wing integrated intensity in order to try to verify the existing observation. + Figure 18. shows that the line wings do show two lobes of emission. with the red lobe being predominantly northeast of the core and the blue lobe," Figure \ref{sfoxx_outflow} shows that the line wings do show two lobes of emission, with the red lobe being predominantly northeast of the core and the blue lobe" +were obtained from the least-squares fitting using the method of Montgomery&ΟDonoghue(1999).,were obtained from the least-squares fitting using the method of \cite{1999DSSN...13...28M}. + Of the 1620 Am stars initially selected. a total of 227 of the total) have been found to pulsate.," Of the 1620 Am stars initially selected, a total of 227 of the total) have been found to pulsate." +" The remaining 1393 stars were deemed as “not found to pulsate"". since low-level pulsation could be present below the SuperWASP detection limits."," The remaining 1393 stars were deemed as “not found to pulsate”, since low-level pulsation could be present below the SuperWASP detection limits." + Table 4. provides a summary of the pulsating Am stars., Table \ref{table:Pulsating Am stars} provides a summary of the pulsating Am stars. + The individual periodograms and phase-folded Hghtcurves are presented in Fig., The individual periodograms and phase-folded lightcurves are presented in Fig. + I., 1. +ΓΙ To place stars on the HR diagram we require values of Tr and log L., To place stars on the HR diagram we require values of $T_\mathrm{eff}$ and $\log L$ . + For stars with wvbyf photometry in the Hauck&Mermilliod(1998) catalogue. we used the code of Moon(1985) to obtain de-reddened indices. and the (5—v. co) grids of Smalley&Kupka(1997) to determine Ty and loge.," For stars with $uvby\beta$ photometry in the \cite{1998A&AS..129..431H} catalogue, we used the code of \citet{1985CommULO..78} to obtain de-reddened indices, and the $b-y$, $c_0$ ) grids of \cite{1997A&A...328..349S} to determine $T_\mathrm{eff}$ and $\log g$." + For stars with only ον photometry the above procedure was used but without the de-reddening step., For stars with only $uvby$ photometry the above procedure was used but without the de-reddening step. + For stars without uvby photometry. Geveva photometry from Rufener(1988) was used and the calibration of Künzlietal.(1997) used to determine Το and logg. assuming zero reddening.," For stars without $uvby$ photometry, Geveva photometry from \cite{1999yCat.2169....0R} was used and the calibration of \cite{1997A&AS..122...51K} used to determine $T_\mathrm{eff}$ and $\log g$, assuming zero reddening." + In all of the above cases. the Torresetal.(2010) relations were usec to determine log£L.," In all of the above cases, the \cite{2010A&ARv..18...67T} relations were used to determine $\log L$." + For stars without suitable band photometry. but with Hipparcos parallaxes (vanLeeuwen. 2007).. spectral energy distributions (SEDs) were constructec using literature broad-band photometry.," For stars without suitable intermediate-band photometry, but with Hipparcos parallaxes \citep{2007A&A...474..653V}, spectral energy distributions (SEDs) were constructed using literature broad-band photometry." + Values of Των wer[7 determined by fitting Kurucz flux distributions to the SEDs and logL determined from the bolometric flux at the earth f£) anc the Hipparcos parallax., Values of $T_\mathrm{eff}$ were determined by fitting Kurucz flux distributions to the SEDs and $\log L$ determined from the bolometric flux at the earth $f_\oplus$ ) and the Hipparcos parallax. + The typical uncertainties are estimatec to be £200 KK in Tey (40.01 in log Τη) and £0.25 in log£L.," The typical uncertainties are estimated to be $\pm 200$ K in $T_\mathrm{eff}$ $\pm 0.01$ in $\log T_\mathrm{eff}$ ) and $\pm 0.25$ in $\log +L$." + The stellar parameters are given in Table 4.., The stellar parameters are given in Table \ref{table:Pulsating Am stars}. + In total around a third of the Am stars investigated have stellar parameters determined., In total around a third of the Am stars investigated have stellar parameters determined. + The sky coverage of the SuperWASP survey overlaps with a large fraction of theKep/er field., The sky coverage of the SuperWASP survey overlaps with a large fraction of the field. + For Am stars with light curves in both theKepler Public archive and the SuperWASP database we have compared the frequencies and amplitudes., For Am stars with light curves in both the Public archive and the SuperWASP database we have compared the frequencies and amplitudes. + This allows us to evaluate the detection limits of SuperWASP., This allows us to evaluate the detection limits of SuperWASP. + Of the 10 stars with both and SuperWASP data. four have clear pulsations with amplitudes2| mmmag(Table 1)). while the other six stars have amplitudes below the SuperWASP detectability limit.," Of the 10 stars with both and SuperWASP data, four have clear pulsations with amplitudes $\ga 1$ mmag (Table \ref{WASP-Kepler-Stars}) ), while the other six stars have amplitudes below the SuperWASP detectability limit." + The analysis (Table 2)) shows good agreement above the nominal SuperWASP mmmag amplitude limit., The analysis (Table \ref{WASP-Kepler-analysis}) ) shows good agreement above the nominal SuperWASP mmag amplitude limit. + There is a suggestion that the amplitudes found using SuperWASP lighteurves are slightly higher than those fromKepler., There is a suggestion that the amplitudes found using SuperWASP lightcurves are slightly higher than those from. + In addition. the SuperWASP frequency can differ from the ‘true’ frequency by a small integer number of dd! aliases.," In addition, the SuperWASP frequency can differ from the `true' frequency by a small integer number of $^{-1}$ aliases." + The comparison also shows that it i5 possible with SuperWASP data to detect frequencies slightly below the mmmag level (Figure 1))., The comparison also shows that it is possible with SuperWASP data to detect frequencies slightly below the mmag level (Figure \ref{WASP-Kepler-pdgram}) ). + Naturally. the variable data quality of ground-based photometry means that not all stars with suitable variability will be detected.," Naturally, the variable data quality of ground-based photometry means that not all stars with suitable variability will be detected." + The pulsating Am stars (see refHR-Diagram)) are concentrated within the fundamental radial mode red and blue edges of Dupretetal.(2005).., The pulsating Am stars (see \\ref{HR-Diagram}) ) are concentrated within the fundamental radial mode red and blue edges of \cite{2005A&A...435..927D}. + This is in agreement with that found by Balonaetal.(2011) for Am stars within theKepler field., This is in agreement with that found by \cite{2011MNRAS.414..792B} for Am stars within the field. + These studies show that pulsating Am stars are concentrated in the cooler region of the instability strip., These studies show that pulsating Am stars are concentrated in the cooler region of the instability strip. + Hot Am stars do not appear to pulsate at the precision of theKepler data., Hot Am stars do not appear to pulsate at the precision of the data. + The standard interpretation of the Am phenomenon is that atomic diffusion — radiative levitation and gravitational settling — in the outer stellar envelope gives rise to the observed atmospheric abundance anomalies., The standard interpretation of the Am phenomenon is that atomic diffusion – radiative levitation and gravitational settling – in the outer stellar envelope gives rise to the observed atmospheric abundance anomalies. + For a typical mid-A star. Tap~8000 K. there are two thin convection zones in the outer envelope.," For a typical mid-A star, $T_{\rm eff}\sim8000$ K, there are two thin convection zones in the outer envelope." + The atmosphere itself is à convection zone a few thousand km thick where ionisation of H drives the convection., The atmosphere itself is a convection zone a few thousand km thick where ionisation of H drives the convection. + Deeper in the atmosphere. at 7~50000 K. the ionisation of also creates a thin convection. zone. where the «-mechanism drives 6 SSct pulsation.," Deeper in the atmosphere, at $T\sim50\,000$ K, the ionisation of also creates a thin convection zone, where the $\kappa$ -mechanism drives $\delta$ Sct pulsation." + It has long been clear that some Am stars and related types do pulsate. particularly the marginal Am stars (labelled spectroscopically as Am: stars). the evolved Am stars (0 DDel orp PPup stars). and some more extreme cases. such as 1188136 (Kurtz1980:: 1981)) and 440765 (Kurtzetal..1995).," It has long been clear that some Am stars and related types do pulsate, particularly the marginal Am stars (labelled spectroscopically as Am: stars), the evolved Am stars $\delta$ Del or $\rho$ Pup stars), and some more extreme cases, such as 188136 \citealt{1980MNRAS.193...29K}; \citealt{1981ApJ...247..969W}) ) and 40765 \citep{1995MNRAS.276..199K}." +. The pulsation modes that we observe in Amstars are low radial order. low spherical degree p modes.," The pulsation modes that we observe in Amstars are low radial order, low spherical degree p modes." + The surface of the star is an anti-node., The surface of the star is an anti-node. + With the low radial order. the vertical wavelength is long compared to the depth of the envelope above the ionisation zone.," With the low radial order, the vertical wavelength is long compared to the depth of the envelope above the ionisation zone." + With the decrease in density with height in the atmosphere. conservation of kinetic energy," With the decrease in density with height in the atmosphere, conservation of kinetic energy" +"simple models, indicates that it is consistent with similar disc dispersal timescales, with differences certainly smaller than the average age difference between CTTs and WTTs.","simple models, indicates that it is consistent with similar disc dispersal timescales, with differences certainly smaller than the average age difference between CTTs and WTTs." + We thank the referee for a constructive report that helped us improve our paper significantly., We thank the referee for a constructive report that helped us improve our paper significantly. + We thank Cathie Clarke for useful comments with regards the statistical analysis., We thank Cathie Clarke for useful comments with regards the statistical analysis. + We thank Aurora Sicilia-Aguilar for providing useful information on the age distribution in the Tr 37 Cluster., We thank Aurora Sicilia-Aguilar for providing useful information on the age distribution in the Tr 37 Cluster. +these gaps are 2.90 and 4.10. respectively.,"these gaps are $2.9\sigma$ and $4.1\sigma$, respectively." + There are 123 cluster members within the region of infrared photometry: Figure 2cc shows the velocity distribution of these galaxies., There are 123 cluster members within the region of infrared photometry; Figure \ref{fig-vhist}c c shows the velocity distribution of these galaxies. + We separate the galaxies into emission (Em) and absorption (Abs) svstems., We separate the galaxies into emission (Em) and absorption (Abs) systems. + Emission ealaxies are (hose objects with a larger Ro value when cross-correlatecl with the emission template: all other galaxies are classified as absorplion., Emission galaxies are those objects with a larger R value when cross-correlated with the emission template; all other galaxies are classified as absorption. + Figure 3. shows (he positions of the 141 cluster members according (ο spectral (vpe., Figure \ref{fig-pos} shows the positions of the 141 cluster members according to spectral type. + There are 18 emission and 123 absorplion svstems., There are 18 emission and 123 absorption systems. + The absorpüon galaxies are more centrally condensed: Che median projected distances of the Abs and Em samples from the ¢D galaxy. are 20.1 and 27'.5. respectively.," The absorption galaxies are more centrally condensed; the median projected distances of the Abs and Em samples from the cD galaxy are $20^{\prime}.1$ and $27^{\prime}.5$, respectively." + This difference is only mareinally statistically significant: a rank-sum test applied to the two samples gives a confidence that the Em and. Abs svstems do not come from the same parent. spatial distiibution., This difference is only marginally statistically significant; a rank-sum test applied to the two samples gives a confidence that the Em and Abs systems do not come from the same parent spatial distribution. + Figure 4. shows the radial velocity distributions of the Abs and Em samples., Figure \ref{fig-veldist} shows the radial velocity distributions of the Abs and Em samples. + The absorption galaxies have cz=14.3634:98 kins ! and o=1095οἱ kms |.," The absorption galaxies have $c\overline{z}=14,363\pm98$ km $^{-1}$ and $\sigma=1095^{+77}_{-64}$ km $^{-1}$." + The dispersion of. the Em sample (o=1123opDGS72) km !) isH comparable.. but. the mean velocityH of. kan ! is lower.," The dispersion of the Em sample $\sigma=1123^{+268}_{-156}$ km $^{-1}$ ) is comparable, but the mean velocity of $c\overline{z}=13,800\pm273$ km $^{-1}$ is lower." + However. a rank-sunm test gives only a confidence that the Em and Abs galaxies follow different. velocitw distributions.," However, a rank-sum test gives only a confidence that the Em and Abs galaxies follow different velocity distributions." +" Dressler&5hectinan(195580) define a test parameter ὁ which measures the deviation ol asubgroups local mean velocity and dispersion. Tora, aud Torus. Hom that of the clusters overall values (7 and σ]: where n is (he number of galaxies in the subgroup."," \citet{dre8b} define a test parameter ${\delta}$ which measures the deviation of asubgroup's local mean velocity and dispersion, $\overline{v}_{local}$ and ${\sigma}_{local}$, from that of the cluster's overall values $\overline{v}$ and $\sigma$ ): where $n$ is the number of galaxies in the subgroup." + The cumulative deviation is δµ. where AN is the number of cluster members and for each galaxy we calculate 9; using the galaxy and its 7—1 nearest neighbors.," The cumulative deviation is $\Delta=\sum^N_{i=1}\delta_i$ , where $N$ is the number of cluster members and for each galaxy we calculate $\delta_i$ using the galaxy and its $n-1$ nearest neighbors." + For a Gaussian velocity distribution with onlv random [Iuctuations. A NX: the presence of substructure can cause A to be significantlo higher than X.," For a Gaussian velocity distribution with only random fluctuations, $\Delta\sim N$ ; the presence of substructure can cause $\Delta$ to be significantly higher than $N$ ." +we better reproduce. the S50 jim counts gives too high CIB Iuctuations.,we better reproduce the 850 $\mu$ m counts gives too high CIB fluctuations. + Future observations with better accuracy will show if these minor discrepancies disappear or are indicative that the bright submm counts are overestimated due to a high. fraction. of eravitationally lensed sources (Perrotta ct al., Future observations with better accuracy will show if these minor discrepancies disappear or are indicative that the bright submm counts are overestimated due to a high fraction of gravitationally lensed sources (Perrotta et al. + 2002). or are simply indicative that. the present phenomenological model is too simple!," 2002), or are simply indicative that the present phenomenological model is too simple!" + The level of the predicted: CLB Uuetuations for dedicated: experiments (as for example. or Alfazina). with respect. to the eirrus confusion noise and. instrumental noise will be discussed in detail in Piat et al. (," The level of the predicted CIB fluctuations for dedicated experiments (as for example or ), with respect to the cirrus confusion noise and instrumental noise will be discussed in detail in Piat et al. (" +in prep).,in prep). + At 170 yam. it is clear from Fig.," At 170 $\mu$ m, it is clear from Fig." + 12. that the recshift distributions of sources that are making the CIB and those making the bulk of the fluctuations are similar., \ref{CIB_fluc_170} that the redshift distributions of sources that are making the CIB and those making the bulk of the fluctuations are similar. + The Iluetuations are not dominated by bright sources just below the detection threshold but by numerous sources at. higher redshift., The fluctuations are not dominated by bright sources just below the detection threshold but by numerous sources at higher redshift. + Thus. in this case. studying the CID fluctuations gives strong constraints on the CLB source population.," Thus, in this case, studying the CIB fluctuations gives strong constraints on the CIB source population." + This is true in the whole submm-mm range., This is true in the whole submm-mm range. + The population of galaxies. where the CLB peaks. will not be accessible by direct. detection in the coming vears.," The population of galaxies, where the CIB peaks, will not be accessible by direct detection in the coming years." + For example.1 will resolve about 20% of the background at 160 sam (Dole et al.," For example, will resolve about $\%$ of the background at 160 $\mu$ m (Dole et al." + 2002). PACS about 50% at 170. jin (Sect.," 2002), about 50 $\%$ at 170 $\mu$ m (Sect." + 7.2.2) andSPIRE less than 104 at 350 jm (Sect., 7.2.2) and less than $\%$ at 350 $\mu$ m (Sect. + 1n conclusion. we have seen that our model gives number counts. redshift distributions. CID intensity and [uctuations that reproduce all the present observations.," In conclusion, we have seen that our model gives number counts, redshift distributions, CIB intensity and fluctuations that reproduce all the present observations." + It can be now used. for future experiment. predictions. in. particular forHoersehel. and. observations.," It can be now used for future experiment predictions, in particular for, and observations." + ForSEPT. a complete and more detailed study. including simulations and a detailed discussion on the confusion. is done in Dole et al. (," For, a complete and more detailed study, including simulations and a detailed discussion on the confusion, is done in Dole et al. (" +2002).,2002). + The confusion is usually defined as the Huctuations of he background sky brightness below which sources cannot »e detected. individuallv., The confusion is usually defined as the fluctuations of the background sky brightness below which sources cannot be detected individually. + Phese fluctuations are caused. by intrinsically discrete. extragalactic sources., These fluctuations are caused by intrinsically discrete extragalactic sources. + In the far-LHi. subnim and mm wavelengths. due to the limited size of the elescopes compared to the wavelength. the confusion noise av an important role in the total noise bugdet.," In the far-IR, submm and mm wavelengths, due to the limited size of the telescopes compared to the wavelength, the confusion noise play an important role in the total noise bugdet." + In fact. he confusion noise is often greater than the instrumental noise. and is thus limiting severely the surveys depth.," In fact, the confusion noise is often greater than the instrumental noise, and is thus limiting severely the surveys depth." + The otal variance a7 of a measurement within a beam cue to oxtragalactic sources with Iuxes less than ορ is given by: where f(0.0) is the two-dimensional bean profile. (in steracians). S the [ux in Jy. and aNdA the cillerential number counts in: 1 |.," The total variance $\sigma^2$ of a measurement within a beam due to extragalactic sources with fluxes less than $_{lim}$ is given by: where $f(\theta, \phi)$ is the two-dimensional beam profile (in steradians), S the flux in Jy, and $\frac{dN}{dS}$ the differential number counts in $^{-1}$ $^{-1}$." +" We call Sj, the confusion The confusion noise. can be determined using two criteria: the so-called photometric and source density criteria (see Dole et al.", We call $_{lim}$ the confusion The confusion noise can be determined using two criteria: the so-called photometric and source density criteria (see Dole et al. + 2002 for a full description)., 2002 for a full description). + Phe photomoetric criterion is related to the quality of the photometry of detected sources. the [Lux measured near ορ being severcly allected by the presence of fainter sources in the beam.," The photometric criterion is related to the quality of the photometry of detected sources, the flux measured near $_{lim}$ being severely affected by the presence of fainter sources in the beam." +" Ht is defined by the implicit equation: where «νο measures the photometric accuracy ancl is usually taken between 3 ancl The source. density criterion is. related to the completness of detected sources above Sy;,,, directly related o the probability to lose sources too close to cach other o be separated.", It is defined by the implicit equation: where $_{phot}$ measures the photometric accuracy and is usually taken between 3 and The source density criterion is related to the completness of detected sources above $S_{lim}$ directly related to the probability to lose sources too close to each other to be separated. +" There. is a threshold. above which the density of sources above ο, is such that à. significant raction of the sources is lost. (it is impossible to separate he individual sources anvmore).", There is a threshold above which the density of sources above $S_{lim}$ is such that a significant fraction of the sources is lost (it is impossible to separate the individual sources anymore). +" For a given source density (with a Poissonian clistribution) N(2S). corresponding to a lux limit ει the probability P to have the nearest source with Hux greater than 8,5, located closer than the distance Ον IS? 6,5, Is the clistance below which sources cannot. be separated and is a function of the beam. profile."," For a given source density (with a Poissonian distribution) $>$ S), corresponding to a flux limit $_{lim}$, the probability P to have the nearest source with flux greater than $_{lim}$ located closer than the distance $\theta_{min}$ is: $\theta_{min}$ is the distance below which sources cannot be separated and is a function of the beam profile." + Lowe note pw. the ENIM. of the beam profile. Gini. can conveniently be expressed. in unitv of py. Ain = kk Ori.," If we note $\theta_{FW}$, the FWHM of the beam profile, $\theta_{min}$ can conveniently be expressed in unity of $\theta_{FW}$, $\theta_{min}$ = $\times$ $\theta_{FW}$." + As an illustration. simulations of source extraction at theAZLPSZSIRTE wavelengths show that k=0.8 should. be achievable (Dole et al.," As an illustration, simulations of source extraction at the wavelengths show that k=0.8 should be achievable (Dole et al." + 2002)., 2002). + Therefore. in the following. we fix k=O0.8.," Therefore, in the following, we fix k=0.8." +" We choose the maximum acceptable probability of not being able to separate the nearest. source D, = 0.1.", We choose the maximum acceptable probability of not being able to separate the nearest source $_{max}$ = 0.1. + In this case. the source density is equal to Lf 17.30 (Table 13) and the corresponding most probable distance is about 1.75Ary. Using Eq. 14.," In this case, the source density is equal to 1 / $\Omega$ (Table \ref{N_omega}) ) and the corresponding most probable distance is about $\times$$\theta_{FW}$ Using Eq. \ref{eq_P}," +" we can derive N(7S) and thus [find the corresponding S,55,,.", we can derive $>$ S) and thus find the corresponding $_{lim}$. + Then. using Eq. 12..," Then, using Eq. \ref{eq_sigma}," + we compute 6., we compute $\sigma$. + The source density criterion leads. to a value of an equivalent Gacusig = NySH., The source density criterion leads to a value of an equivalent $_{density}$ = $\frac{S_{lim}}{\sigma}$. + LW Ἐν is greater than standard. values of quos; (8 to 5). then the confusion noise is given by the source density. criterion.," If $_{density}$ is greater than standard values of $_{phot}$ (3 to 5), then the confusion noise is given by the source density criterion." + Lf not. then the photometrie criterion has to be used to derive the confusion noise.," If not, then the photometric criterion has to be used to derive the confusion noise." + The classical confusion limitof 1 source per 30 beams corresponds to k=l and P-0.1., The classical confusion limitof 1 source per 30 beams corresponds to k=1 and $\sim$ 0.1. + Nevertheless it can still leac to mediocre photometry for very steep The transition between the photometric and. source density criterion is around 200 yaa. depending on telescope," Nevertheless it can still lead to mediocre photometry for very steep The transition between the photometric and source density criterion is around 200 $\mu$ m, depending on telescope" +of the IT aud I. band images taken with Sofl. The kuown foreground stars aud embedded cloud members (see C1999 and C2000) were nof used in computing this map.,of the H and $_{\rm s}$ band images taken with SofI. The known foreground stars and embedded cloud members (see C1999 and C2000) were not used in computing this map. +" Because there were πο observations of the wnreddened backeround stellar feld (uecessary for the determination of the zero point of the extinction scale). we adopted a «II EK,> color for this background."," Because there were no observations of the unreddened background stellar field (necessary for the determination of the zero point of the extinction scale), we adopted a $<$ $-$ $_{\rm s}>$ color for this background." + Roeeurdiung this choice. Alves et al. (," Regarding this choice, Alves et al. (" +1998) found =0.20+£0.13 mae towards a complicated region iu Cyenus. essenutiallv at the Calactie plane. while Alves et al. (,"1998) found $<$ $-$ $_{\rm s}> = 0.20 \pm 0.13$ mag towards a complicated region in Cygnus, essentially at the Galactic plane, while Alves et al. (" +2001) found «II Ky>=0.12+0.08 nag owards a line-of-seht towards the Galactic bulge. at b=~7°.,"2001) found $<$ $-$ $_{\rm s}> = 0.12\pm0.08$ mag towards a line-of-sight towards the Galactic bulge, at $b \simeq 7^{\circ}$." + The stellar vackeround towards the Chamaeleon complex is probably even better behaved. because it lies at b215°.," The stellar background towards the Chamaeleon complex is probably even better behaved, because it lies at $b \simeq 15^{\circ}$." +" On the other haud. the «II KR, color of the North Galactic xole from 2MASS data is also ~0.12 mag (AL Lombardi. oivate conumuimuication). so we decided. to adopt this as he mean Chamacleon background color."," On the other hand, the $<$ $-$ $_{\rm s}$$>$ color of the North Galactic pole from 2MASS data is also $\sim 0.12$ mag (M. Lombardi, private communication), so we decided to adopt this as the mean Chamaeleon background color." + We also adopt a conservative intrinsic dispersion iu this nieasureimoent of 12 mag. ie. an mtrimsic extinction measurement rnis of Ay=2 mug.," We also adopt a conservative intrinsic dispersion in this measurement of 0.13 mag, i.e. an intrinsic extinction measurement rms of $_{\rm V} \simeq 2$ mag." + We show our map in Fie., We show our map in Fig. + 7 for the area around Cha πα 1. 5. S. 10. and LL. where companion caucdidates were detected.," 7 for the area around Cha $\alpha$ 4, 5, 8, 10, and 11, where companion candidates were detected." + The spatial resolution of the map is 50 and the extinction Ay ranges from 1 to 22 mae., The spatial resolution of the map is $^{\prime \prime}$ and the extinction $_{\rm V}$ ranges from 4 to 22 mag. + Each position of this map is associated with a total extinction which icludes the whole Cha I dark cloud aud. its fore- aud οackeround., Each position of this map is associated with a total extinction which includes the whole Cha I dark cloud and its fore- and background. + C2000 derived the individual reddening by comparing the observed colors to iutrisic colors for the known spectral types., C2000 derived the individual reddening by comparing the observed colors to intrinsic colors for the known spectral types. + Those absorptions range from Ayz0.0 to 1.0 mae (C2000) and are simaller than the extinction values towards their imuuediate background. as measured oe1 the extinction map.," Those absorptions range from $_{\rm V} \simeq 0.0$ to 1.0 mag (C2000) and are smaller than the extinction values towards their immediate background, as measured in the extinction map." + The total extinction towards the oemediate backeround of the primarics shown in Fie., The total extinction towards the immediate background of the primaries shown in Fig. + 7 range from Ayz 6 to 8 mag., 7 range from $_{\rm V}\simeq $ 6 to 8 mag. +" Ποσο, the primaries are most certainly located near the frout edge of the dark cloud."," Hence, the primaries are most certainly located near the front edge of the dark cloud." + Cha Ta 5 is an M6 dwarf with Πα and Nav chussion as well as Lithia absorption: no radial velocity variations were detected in UVES spectra. consistent with membership to the Cha I association (C1999. C2000: Joerseus Carenther 2001).," Cha $\alpha$ 5 is an M6 dwarf with $\alpha$ and X-ray emission as well as Lithium absorption; no radial velocity variations were detected in UVES spectra, consistent with membership to the Cha I association (C1999, C2000; Joergens Guenther 2001)." + Based ou a comparison of, Based on a comparison of +brightness profile for the remaiming galaxies.,brightness profile for the remaining galaxies. + With ouly 3 exceptions these profiles are well reproduced by a Nuker law fit., With only 3 exceptions these profiles are well reproduced by a Nuker law fit. + We are then able to separate carly-type galaxies on the basis of the slope of their nuclear brightuess profiles (>). iuto core (5x: 0.3) and power-law (+> 0.5) galaxies: we also found 3 galaxies with an intermediate slope (0.3< 0.5).," We are then able to separate early-type galaxies on the basis of the slope of their nuclear brightness profiles $\gamma$ ), into core $\gamma \leq 0.3$ ) and power-law $\gamma \geq 0.5$ ) galaxies; we also found 3 galaxies with an intermediate slope $ 0.3 < \gamma < 0.5$ )." + We preferred this classification with respect to the traditiona morphological scheme G.c. E and SO ealaxies) since it is well kuown that it is often difficult to unuiuuabieuouslv distineuish between these two classes of objects (vandeuBergh2001)., We preferred this classification with respect to the traditional morphological scheme (i.e. E and S0 galaxies) since it is well known that it is often difficult to unambiguously distinguish between these two classes of objects \citep{vandenbergh04}. +. We recovered the already known difference iu the luuimositv between core aud power-law ealaxics. since the most luminous galaxies are exchisively core galaxies.," We recovered the already known difference in the luminosity between core and power-law galaxies, since the most luminous galaxies are exclusively core galaxies." + Below Ay~212. however. the wo classes coexist and core galaxies are associated to objects as faint as Mi~2ὃν," Below $M_K \sim -24.2$, however, the two classes coexist and core galaxies are associated to objects as faint as $M_K \sim -21.8$." +" A well defined correlation between the radius and he surface brielitucss at the profile break. ry, and gy. the vcore fundamental plane” .ds present also in our süuple and it is ndistinguishable from that secu in the other samples. purely optically selected. studied."," A well defined correlation between the radius and the surface brightness at the profile break, $r_b$ and $\mu_b$, the “core fundamental plane”, is present also in our sample and it is indistinguishable from that seen in the other samples, purely optically selected, studied." +" ""Thus. our sample does not reflec any significant difference in the host galaxies with respect to the normal galaxies population."," Thus, our sample does not reflect any significant difference in the host galaxies with respect to the normal galaxies population." +" Concerning the relationship between hosts magnitude and radio-power. we found that onlv core-galaxies are radio-euütters at level larger than £,>2.51072 W |. confirniug the sugecstiousOO bv De Ruiter et al. ("," Concerning the relationship between host's magnitude and radio-power, we found that only core-galaxies are radio-emitters at level larger than $L_r > 2.5 \times 10^{21}$ W $^{-1}$, confirming the suggestions by De Ruiter et al. (" +2005).,2005). + But below this threshold the two populatious of galaxies caunot be readily differentiated., But below this threshold the two populations of galaxies cannot be readily differentiated. +" Siuce no strong radio-source (L,>>1072 W 1) is associated to a host with Mg>2| regardless of the profile class. we are loft with the ambiguity on the origin of this threshold in racdio-huninosity for power-law galaxies: it can be due voth to a different nuclear structure of the two classes or to a difference iu he host ealaxvs bhDuninositv. since sower-law ealaxies ouly scarcely populate the Heh cud of he optical uumnositv distributions. where the brightest racdio-sources are found."," Since no strong radio-source $L_r > 10^{22}$ W $^{-1}$ ) is associated to a host with $M_K > -24$ regardless of the profile class, we are left with the ambiguity on the origin of this threshold in radio-luminosity for power-law galaxies: it can be due both to a different nuclear structure of the two classes or to a difference in the host galaxy's luminosity, since power-law galaxies only scarcely populate the high end of the optical luminosity distributions, where the brightest radio-sources are found." + The imost significantzl difference concerningoO the optical xoperties of our sample with respect to previous studies is the higher fraction of nucleated galaxies. Sce we detected an optical (or infrared) nucleus in —6! of the objects.," The most significant difference concerning the optical properties of our sample with respect to previous studies is the higher fraction of nucleated galaxies, since we detected an optical (or infrared) nucleus in $\sim 65\%$ of the objects." + This is not at all unexpected. since our sample is deliberately biased to favour the inclusion of active ealaxies.," This is not at all unexpected, since our sample is deliberately biased to favour the inclusion of active galaxies." + Addressing the multmwaveleugth properties of these nuclei will be the aim of the two forthcoming papers of the series., Addressing the multiwavelength properties of these nuclei will be the aim of the two forthcoming papers of the series. +critical value maj. which can be as low as 1. depending on the shape of the curent channel ol the fIux rope (e.g.DemoulinandAulanier2010).,"critical value $n_{\rm crit}$, which can be as low as $1$, depending on the shape of the current channel of the flux rope \citep[e.g.][]{demoulin_aulanier2010}." +. For a 3D anchored flux rope. as is the case here. it is diffieult to obtain an analvtical determination of n4 lor the instability or loss of equilibrium of the flix rope (IsenbergandForbes2007).," For a 3D anchored flux rope, as is the case here, it is difficult to obtain an analytical determination of $n_{\rm crit}$ for the instability or loss of equilibrium of the flux rope \citep{isen_forbes2007}." +. The exact critical point for the onset of the torus instability would depend on the detailed 3D magnetic fied configuration., The exact critical point for the onset of the torus instability would depend on the detailed 3D magnetic field configuration. + On (he other hand. a substantial amount of twist has been transported into the corona al the onset of eruption.," On the other hand, a substantial amount of twist has been transported into the corona at the onset of eruption." +" At /=2.5. the sell-helieitv of the emerged[τικ rope reaches about =1.020%ne where 4, 15 the total magnetic [flux in the rope. corresponding to field lines in the [Iux rope winding about the central axis by about 1.02 rotations between the auchored foot points."," At $t=2.5$, the self-helicity of the emergedflux rope reaches about $-1.02 \Phi_{\rm rope}^2$, where $\Phi_{\rm rope}$ is the total magnetic flux in the rope, corresponding to field lines in the flux rope winding about the central axis by about 1.02 rotations between the anchored foot points." + This suggests the possible development of the helical kink instability of the flux rope (e.g.lloodandPriest1981:TorokWien:2005:FanGibson 2007).," This suggests the possible development of the helical kink instability of the flux rope \citep[e.g.][]{hood_priest1981, toeroek_kliem2005, fan_gibson2007}." +. The erupting flux rope is found (to undergo substantial writhing or kinking motion as can be seen in the sequences of images (also the movies in the electronic version) in Figures 4. and 5.., The erupting flux rope is found to undergo substantial writhing or kinking motion as can be seen in the sequences of images (also the movies in the electronic version) in Figures \ref{fig4} and \ref{fig5}. + We also find that the trajectory for the eruption of the Πας rope is not radial because of the ambient coronal magnetic Ποια: the erupting flix rope is deflected southward ancl eastward from the local radial direction (see Figures 4.. 5.. and 7 and the associated movies).," We also find that the trajectory for the eruption of the flux rope is not radial because of the ambient coronal magnetic field: the erupting flux rope is deflected southward and eastward from the local radial direction (see Figures \ref{fig4}, \ref{fig5}, and \ref{fig7} and the associated movies)." + Using the apex location of the erupting flux rope cavity al /=3.25 (Figure 7)). we lind that the eruptng trajectory at (hat time is deflected by. 2.37 southward and 1.37 eastwarel from the local radial direction at the center of flux emergence. and further deflection of the trajectory continues with time.," Using the apex location of the erupting flux rope cavity at $t=3.25$ (Figure \ref{fig7}) ), we find that the erupting trajectory at that time is deflected by $2.3^{\circ}$ southward and $1.3^{\circ}$ eastward from the local radial direction at the center of flux emergence, and further deflection of the trajectory continues with time." + Since the local radial direction at the center of the flux emergence corresponds to 1.175 and 24 W from the solar disk center (or the Iine-ol-sight. the deflection during the eruption is sending the [αν rope towards the line-ol-sight in the east-west direction. but further soutliward away. [rom the line-ol-sight in the north-south direction.," Since the local radial direction at the center of the flux emergence corresponds to $7.1^{\circ}$ S and $24^{\circ}$ W from the solar disk center (or the line-of-sight), the deflection during the eruption is sending the flux rope towards the line-of-sight in the east-west direction, but further southward away from the line-of-sight in the north-south direction." + This is consistent with the observed halo of the CME seen in LASCO C? and C3 coronagraphs (Figure 2 in WKataokaetal.(2009) ancl Figure 1 in Ravinclra (2010))). where the north-south and east-west asviuumetries of the halo distribution indicate that the direction of ejection is more southward and less westward than what would have been expected for a radial ejection from the location of the source region on (the solar disk.," This is consistent with the observed halo of the CME seen in LASCO C2 and C3 coronagraphs (Figure 2 in \citet{kataokaetal2009} and Figure 1 in \citet{ravindra_howard2010}) ), where the north-south and east-west asymmetries of the halo distribution indicate that the direction of ejection is more southward and less westward than what would have been expected for a radial ejection from the location of the source region on the solar disk." + Figure & shows the coronal magnetic field as viewed from (he side (panels a and b) aud viewed Irom the observing perspective (panels c aud d) just before the onset of eruption al |—2.45. compared with the Iinode XRT image of the region (panel e) just before the flare.," Figure \ref{fig8} shows the coronal magnetic field as viewed from the side (panels a and b) and viewed from the observing perspective (panels c and d) just before the onset of eruption at $t=2.45$, compared with the Hinode XRT image of the region (panel e) just before the flare." + We see that the morphology of the coronal magnetic field ancl its connectivity are very similar io those shown in the N-rav image., We see that the morphology of the coronal magnetic field and its connectivity are very similar to those shown in the X-ray image. + To understand the nature of the bright X-ray. sigmoid in(he image. we have iclentified the reeion of significant magnetic energv dissipation ancl heating in the simulated magnetic field using both the electric current density J=|VxB| and the increase of entropy AS=C.N In(p/p?).," To understand the nature of the bright X-ray sigmoid inthe image, we have identified the region of significant magnetic energy dissipation and heating in the simulated magnetic field using both the electric current density $J \equiv | \nabla \times {\bf B} |$ and the increase of entropy $\Delta S = C_v \Delta \ln (p/{\rho}^{\gamma} )$ ." + As pointed out in Section 2.. since we are solving the total οποιον equation in conservative form. numerical dissipation of magnetic," As pointed out in Section \ref{sec:model}, , since we are solving the total energy equation in conservative form, numerical dissipation of magnetic" +"systems within the sub-mJy population is higher than that in local galaxy samples and (i) the pairing properties of the sub-mJyv radio sample (with the given biases such as resolution ellects) are not significantly cülferent from. those of ""Ποια galaxies at the same magnitude limit.",systems within the sub-mJy population is higher than that in local galaxy samples and (ii) the pairing properties of the sub-mJy radio sample (with the given biases such as resolution effects) are not significantly different from those of `field' galaxies at the same magnitude limit. + During part of this study AC was supported by the State Scholarship Foundation of Cireece. (XY)., During part of this study AG was supported by the State Scholarship Foundation of Greece (IKY). + We thank the referee for helpful comments and suggestions that improved this paper., We thank the referee for helpful comments and suggestions that improved this paper. +(sec Tovmassian. Yam TViersch 2002: Plionis VPovmassian 2004).,"(see Tovmassian, Yam Tiersch 2002; Plionis Tovmassian 2004)." + AIP acknowledges ME by the Mexican Government erant No., MP acknowledges funding by the Mexican Government grant No. + CONAC\T-201-39679., CONACyT-2002-C01-39679. +" SB acknowledges the hospitality of INAOL ""m this ""m was completed.", SB acknowledges the hospitality of INAOE where this work was completed. +deviations ofz0.2 as shown in Fig. 7..,deviations of$\approx 0.2$ as shown in Fig. \ref{fig:circ_rperi_dist_evol}. +" Thus, as with mass dependence, the redshift evolution of the mean/median out to z—5 is comparable to the distribution width itself."," Thus, as with mass dependence, the redshift evolution of the mean/median out to $z=5$ is comparable to the distribution width itself." + What drives the redshift evolution of satellite orbits?, What drives the redshift evolution of satellite orbits? + One possibility is that redshift dependence is simply a manifestation of the trends with mass from the previous section., One possibility is that redshift dependence is simply a manifestation of the trends with mass from the previous section. +" Analytical triaxial collapse models of halo formation (???) predict a self-similarity in the nature of matter infall with redshift at fixed M/M.(z), since it is at M. "," Analytical triaxial collapse models of halo formation \citep{BBKS86,EisLoe95,SheMoTor01} predict a self-similarity in the nature of matter infall with redshift at fixed $M/M_*(z)$ , since it is at $M_*$ " + 2007:Swainctal.2008).. (Charbo," \citep{charb02,tin07,swain08}. \citep{knutson07b,charb08}," +nneanetal.2005:—Deming2007) (auutsonetal.," \citep{charb05,dem05,har07} \citep{knutson07}." +2007).. Po. S ?7 Charbonneanetal.(2008))) Burrowsetal.(2006).. Fortuevetal.(2006)x/ Burrowsetal.(2007a).. (Werneretal.(2001))). (Charbouneauetal.2005).. (Demingctal.2005:Ikuutsou2008).. (Ixuutsonctal.2007).. (Tarrinetou2007) (Demingetal.2007).," \cite{torres08} $_{orb}$ $\lesssim$ $^{-3}$ \cite{charb08}) \citet{burr06}, \citet{fort06} \citet{burr07}, \citet{werner04}) \citep{charb05}, \citep{dem05,knutson07b}, \citep{knutson07}, \citep{har07} \citep{dem07}." +. addition low resolution spectra of 2 transiting plaucts were obtained with the IRS spectrometer between ~ 7 and 15 4500: (Coillingziretal.2007) anc (Richardsonctal.2007)., addition low resolution spectra of 2 transiting planets were obtained with the IRS spectrometer between $\sim$ 7 and 15 $\mu m$: \citep{grill07} and \citep{rich07}. +". Recently ai detection of an atimospheric feature attributed to water has heen claimed by Tinettietal.(2007) and Baran(2007) by studving the transit flux ratios of aud158b.. respectively,"," Recently a detection of an atmospheric feature attributed to water has been claimed by \citet{tin07} and \citet{barman07} by studying the transit flux ratios of and, respectively." + Burrowsctal.(2007a) analyzed the secoudazy transiuission spectra of at all 1 infrared IRAC Spitzer channels observed by Kuutsonetal.(2008) aud sugeested the observations are consistcut with an atimospheric thermal inversion laver aud vet ΠΠμοι stratospheric absorber., \citet{burr07} analyzed the secondary transmission spectra of at all 4 infrared IRAC Spitzer channels observed by \citet{knutson07b} and suggested the observations are consistent with an atmospheric thermal inversion layer and yet unknown stratospheric absorber. + A detailed study of the IR secondary eclipse plauctary spectra of05...189733b..TrES-1.. and uou-eclipsiug179919b.. aud bv Burrowsetal.(20075) sugeests that the presenceof such a stratospheric absorber might be dependent on the flux from the star at the sub-stellar point on the planet as well as second order effects like metallicity and planetary surface evavity.," A detailed study of the IR secondary eclipse planetary spectra of, and non-eclipsing, and by \citet{burr07b} suggests that the presenceof such a stratospheric absorber might be dependent on the flux from the star at the sub-stellar point on the planet as well as second order effects like metallicity and planetary surface gravity." + In the Burrowsotal.(2007b) interpretation auets with ligh sub-stellar point flux (e.g. 105υ.. OGLE-Ti-56b. OCLE-Tr-132b. anc NO-3))) would have a stratospheric laver aud a water cative in ciuission while planets with lower fluxcs(NO-1b..TYES-1.. and 189733h)) woul wave no such laver and a water feature iu absorption.," In the \citet{burr07b} interpretation planets with high sub-stellar point flux (e.g., , OGLE-Tr-56b, OGLE-Tr-132b, and ) would have a stratospheric layer and a water feature in emission while planets with lower fluxes, and ) would have no such layer and a water feature in absorption." + Fortuevetal.(2007) also suggest a simular division of auetarv spectra based on Πιοοι! stella flux., \citet{fort07b} also suggest a similar division of planetary spectra based on incident stellar flux. + Bascc on the planetary sub-stelar point flux from the star. voth Burrowsetal.(2007) and Fortueyetal.(2007) predict that should not exhibit a therm version m its atimosphere.," Based on the planetary sub-stellar point flux from the star, both \citet{burr07b} and \citet{fort07b} predict that should not exhibit a thermal inversion in its atmosphere." + We present observations ofthe infrared spectral energv distribution (SED) ofthe planet (AIcCulough in all | TRAC channels obtained during secondary eclipses with the TRAC of Telescope., We present observations of the infrared spectral energy distribution (SED) ofthe planet \citep{xo1} in all 4 IRAC channels obtained during secondary eclipses with the IRAC of . +.. Bycomparing our ~ [8 gin SED with atmospheric models. we test for the presence of a thermal inversion laver in XO-Ib..," Bycomparing our $\sim$ 4-8 $\mu m$ SED with atmospheric models, we test for the presence of a thermal inversion layer in ." +of view and fitting results.,of view and fitting results. +" To help making decisions in this situation, we need to know what is the best achievable fitting result from various models or parametrizations with the same number of parameters."," To help making decisions in this situation, we need to know what is the best achievable fitting result from various models or parametrizations with the same number of parameters." + This will serve as a fiducial criteria for us to choose a model., This will serve as a fiducial criteria for us to choose a model. + The figure of merit introduced in this paper is to help to define what is the best., The figure of merit introduced in this paper is to help to define what is the best. +" As is well known, besides SNIa observations, there exist lots of other experiments probing different aspects of dark energy and we will have many more data from these experiments (Albrechtetal.2006)."," As is well known, besides SNIa observations, there exist lots of other experiments probing different aspects of dark energy and we will have many more data from these experiments \citep{DETF1}." +". However, in terms of constraining the evolution of w(z), SNIa approach is the most sensitive and direct one."," However, in terms of constraining the evolution of $w(z)$, SNIa approach is the most sensitive and direct one." +" Other methods, such as CMB and cluster counts, are primarily good for the energy density constraint."," Other methods, such as CMB and cluster counts, are primarily good for the energy density constraint." + But it will be advantageous to test all the parametrization with all the combined data sets in the future., But it will be advantageous to test all the parametrization with all the combined data sets in the future. + The current analysis in this paper could be directly generalized to the case with multi-experiments by maximizing the product of the likelihood of each experiment., The current analysis in this paper could be directly generalized to the case with multi-experiments by maximizing the product of the likelihood of each experiment. + It is worth noting that the best parametrization of dark energy models for SNIa data may not necessarily be the best one for other observational data., It is worth noting that the best parametrization of dark energy models for SNIa data may not necessarily be the best one for other observational data. + We will report that in a preparing work., We will report that in a preparing work. + This work is supported by National Natural Science Foundation of China under Grant No., This work is supported by National Natural Science Foundation of China under Grant No. + 10473007 and No., 10473007 and No. + 10503002 and Shanghai Commission of Science and technology under Grant No., 10503002 and Shanghai Commission of Science and technology under Grant No. + 06QA14039., 06QA14039. + Enussion frou Πο aasers has been found iu a few ealaxies. exhibiting appareut isotropic Lluiinositfies a milliou. tines higher than iu typical stellar miasers (DosTaschick&Baan1955:Braatz.Wilson.Παικοα199 1).," Emission from $_2$ O masers has been found in a few galaxies, exhibiting apparent isotropic luminosities a million times higher than in typical stellar masers \cite{DosSantosLepine1979,GardnerWhiteoak1982,ClaussenHeiligmanLo1984,HenkelGuestenWilson1984,HaschickBaan1985,BraatzWilsonHenkel1994}." + The detection rate of these imeeamascrs is ναν low. ic. about amone Sevtert ealaxies (Draatz.Wilson.&1997) aud almost zero amoung radio galaxies (0.9.. Henkelal.1998)).," The detection rate of these megamasers is very low, i.e. about among Seyfert galaxies \cite{BraatzWilsonHenkel1997} and almost zero among radio galaxies (e.g., \citeNP{HenkelWangFalcke1998}) )." + The maser is associated with deuse and warn nateral. possibly a molecular torus or disk. around an active ealactie nucleus (AGN).," The maser is associated with dense and warm material, possibly a molecular torus or disk, around an active galactic nucleus (AGN)." + The ACN apparcutly produces the seed radio photons aud the X-ray photous needed to pump the masing material (Neufeld.Maloney.&Conger 1991)., The AGN apparently produces the seed radio photons and the X-ray photons needed to pump the masing material \cite{NeufeldMaloneyConger1994}. + With the he» of Verv Long Bascline Tuterterometiy (VLDI) inceaimasers can be used to investigate the smallscale structure of an AGN in erea detail., With the help of Very Long Baseline Interferometry (VLBI) megamasers can be used to investigate the small-scale structure of an AGN in great detail. + In the case of NGC 1258 this as helped to establish the preseuce of a thin. warped disk around the nucleus. to determine the black hole mass. aud even to measure the precise distauce to this ealaxy (Mivoshietal.1995:IHerrustein1999).," In the case of NGC 4258 this has helped to establish the presence of a thin, warped disk around the nucleus, to determine the black hole mass, and even to measure the precise distance to this galaxy \cite{MiyoshiMoranHerrnstein1995,HerrnsteinMoranGreenhill1999}." + Finding new ueganuaser galaxies is therefore of prime interest., Finding new megamaser galaxies is therefore of prime interest. + The only cler trend hat has emerged iu recent vears is that meeamasers are exclusively found in type 2 AGN. Le. those Sevterts aud LINERs which are expected to be obscured bv a molecular torus according to the unified scheme (Antonucci1993).," The only clear trend that has emerged in recent years is that megamasers are exclusively found in type 2 AGN, i.e. those Seyferts and LINERs which are expected to be obscured by a molecular torus according to the unified scheme \cite{Antonucci1993}." +". Muy have high absorbing colhuun depths interred frou, X-ray spectroscopy,", Many have high absorbing column depths inferred from X-ray spectroscopy. + There is also an indication of an excess of mieezinasers m highly inchned galaxies (Draatz.Wilsou.&Henkel1997:Falckeetal. 2000).," There is also an indication of an excess of megamasers in highly inclined galaxies \cite{BraatzWilsonHenkel1997,FalckeWilsonHenkel2000}." +. Tere we report the discovery of a hitherto uudoetected arc very huuinuous mcewmaser m the Sevfert ealaxy Mrk 248 dunug a radio flare of the ACN., Here we report the discovery of a hitherto undetected and very luminous megamaser in the Seyfert galaxy Mrk 348 during a radio flare of the AGN. + Ak 3148 (NGC 262. 2=0.01503 Tnchra.Vogeleyv.&Geller 1999.. huninosity distance D=62.5 AIpe for - converted into the Calactic Standard of Rest aud fy=15 laus ty). is a welkstucdied Seyfert 2 galaxv with broad cluissiou-lines in polarized light (Miller& 1990).," Mrk 348 (NGC 262, $z=0.01503$ \citeNP{HuchraVogeleyGeller1999}, luminosity distance $D=62.5$ Mpc for $z$ converted into the Galactic Standard of Rest and $H_{0}=75$ km $^{-1}$ $^{-1}$ ), is a well-studied Seyfert 2 galaxy with broad emission-lines in polarized light \cite{MillerGoodrich1990}." +. The galaxy is classified as au SO with a rather low inclination (7=167. see Braatz.Wilson.&Ieukel1997 )).," The galaxy is classified as an S0 with a rather low inclination $i=16^\circ$, see \citeNP{BraatzWilsonHenkel1997}) )." + Ground-based (Simpsonctal.1996) and IIubble Space Telescope inagiug (Falcke.Wilson.&Simpson1998) jnve revealed a dust lane crossing the uucleus aud au excitation cone iu enüssion-liues., Ground-based \cite{SimpsonMulchaeyWilson1996} and Hubble Space Telescope imaging \cite{FalckeWilsonSimpson1998} have revealed a dust lane crossing the nucleus and an excitation cone in emission-lines. + Cünga observations found lard A-rayv cluission and a hieh absorbing cohuun depth of, Ginga observations found hard X-ray emission and a high absorbing column depth of +For AGN (riggering in a gas rich galaxy. the occurrence of a tidal perturbation may be more relevant than its duration (Ixeel1996).,"For AGN triggering in a gas rich galaxy, the occurrence of a tidal perturbation may be more relevant than its duration \citep{keel96}." +. An hyperbolic encounter may well trigger a racial flow in the innermost regions of a gas-rich galaxy., An hyperbolic encounter may well trigger a radial flow in the innermost regions of a gas-rich galaxy. + The time needed by (he companion to move away by 30 Άρο is ~1.x107dauNeA0041ms; vr.," The time needed by the companion to move away by 30 Kpc is $\sim 1.\times +10^8 d_{30Kpc} \Delta v^{-1}_{300 Km s^{-1}}$ yr." + The (mescale for a chump of eas to fall from the outer regions of the nucleus (a few hundreds of parsecs) to the inner central pc is Z0.1 Gyr (Dekki(2000))). and this can be considered a lower limit to the lime needed for the onset of the active nucleus.," The timescale for a clump of gas to fall from the outer regions of the nucleus (a few hundreds of parsecs) to the inner central pc is $\simgt +0.1$ Gyr \citet{b00}) ), and this can be considered a lower limit to the time needed for the onset of the active nucleus." +" Therefore. an hyperbolic encounter with moderate Ar, ὅσα be such that the companion escape from the close vicinity (260 KXpe) of the Sevlert galaxy. leaving a non-interacting Sevlert 1 nucleus."," Therefore, an hyperbolic encounter with moderate $\Delta v_r$ can be such that the companion escape from the close vicinity $\approx 60 $ Kpc) of the Seyfert galaxy, leaving a non-interacting Seyfert 1 nucleus." + If obsceuration is significant. or if the AGN power is small (because of low accretion rate. or of an undermassive central black hole). then a longer timescale may be necessary before a Sevfert 1 nucleus is actually detected.," If obscuration is significant, or if the AGN power is small (because of low accretion rate, or of an undermassive central black hole), then a longer timescale may be necessary before a Seyfert 1 nucleus is actually detected." + While BIRGs and Sv2 galaxies have richer environments than Svls at distances X60Ape. the cumulative distribution of the projected separation for the first companion (Fig. 3))," While BIRGs and Sy2 galaxies have richer environments than Sy1s at distances $\simlt 60 Kpc$, the cumulative distribution of the projected separation for the first companion (Fig. \ref{fig03}) )" + shows that the environmental difference for Svls. with respect to Syv2s and DIBGs decreases dramatically bevond zz 120 pce.," shows that the environmental difference for Sy1s, with respect to Sy2s and BIRGs decreases dramatically beyond $\approx$ 120 Kpc." +" This means that. while Sv2s and. BIRGs have close companions. Svls do have companions. but at higher distances (dj,ZLOO Ixpe)."," This means that, while Sy2s and BIRGs have close companions, Sy1s do have companions, but at higher distances $d_p \simgt 100$ Kpc)." + 5v1s do not show close companions simply because any. activitv-triggering interaction (ook place in (he past. aud. on average. Svl galaxies would not be considered interacting following our statistical criterions.," Sy1s do not show close companions simply because any activity-triggering interaction took place in the past, and, on average, Sy1 galaxies would not be considered interacting following our statistical criterions." + The limitations of our analysis regarding small Companion galaxies (De&104v pe) leave open other main possibility to account for tvpe-I activity., The limitations of our analysis regarding small companion galaxies $D_C \simlt 10 Kpc$ ) leave open other main possibility to account for type-1 activity. + It has been proposed that Sv1s mav be the result of a “minor merger which purportedly may not lead to a dramatic star formation close to the center of the galaxies and hence to heavy obscuration (deRobertiV.(1993). Daniguchi (1999))).," It has been proposed that Sy1s may be the result of a “minor merger"" which purportedly may not lead to a dramatic star formation close to the center of the galaxies and hence to heavy obscuration \citet{dr98}, \citet{t99}) )." + N-body simulations of minor mergers show that they produce disturbances in (he morphology of the larger galaxy in the first Gyr of the onset of the merger. but do not destrov the galactic disk (Walker. (1996))).," N-body simulations of minor mergers show that they produce disturbances in the morphology of the larger galaxy in the first Gyr of the onset of the merger, but do not destroy the galactic disk \citet{w96}) )." + Corbin(2000). did not found higher levels of asxaimetry in Sevlert galaxies (han in normal galaxies (in agreement with our work. he found that the most asvimmetric galaxies were interacting svstems with ILHIL-like spectra).," \citet{c00} did not found higher levels of asymmetry in Seyfert galaxies than in normal galaxies (in agreement with our work, he found that the most asymmetric galaxies were interacting systems with HII-like spectra)." + He concluded (hat. if minor mergers (rigger AGNs. they appear to do so only in the late stages of the mergers ( 1 Gyr after the merger onset).," He concluded that, if minor mergers trigger AGNs, they appear to do so only in the late stages of the mergers $\sim$ 1 Gyr after the merger onset)." + Minor mergers also boost the star formation of the larger galaxy. but this process is nol necessarily very dramatic (the induced SER max be as low as 21./ yr) especially alter the first 0.5 Gvr (Rudnick.Rix.&Ixennieutt. (2000))).," Minor mergers also boost the star formation of the larger galaxy, but this process is not necessarily very dramatic (the induced SFR may be as low as $\approx 2 M_{\odot}/yr$ ) especially after the first $0.5$ Gyr \citet{r00}) )." + The previous mechanisms suggest a revision (hat complement the unification scheme [or Sevfert galaxies. and favor the idea of a long timescale to let type-l AGN emerge.," The previous mechanisms suggest a revision that complement the unification scheme for Seyfert galaxies, and favor the idea of a long timescale to let type-1 AGN emerge." + It is interesting to stress that times for the onset of this kind of activity are in agreement with (he time needed to let any unbound companion [lv at least few tens of Ixpce. or to have a full," It is interesting to stress that times for the onset of this kind of activity are in agreement with the time needed to let any unbound companion fly at least few tens of Kpc, or to have a full" +Archival oobservations performed in 2003 anc 2009 reveal quiescent spectra that fit to a neutron star atmosphere model with KI~72 eV. viekding a thermal bolometric luminosity of Ly~GLO(D/5.5kpc)?ergs+ (Degenaar&Wijnancds 2011).,"Archival observations performed in 2003 and 2009 reveal quiescent spectra that fit to a neutron star atmosphere model with $kT^{\infty} \sim 72$ eV, yielding a thermal bolometric luminosity of $L_q \sim 6 \times 10^{32}~(D/5.5~\mathrm{kpc})^2~\lum$ \citep{deeg_wijn2011}." +. Our new oobservation demonstrates that within two months alter he cessation of the recent accretion outburst. the thermal lux and neutron star temperature are elevated. above the quicscent base level measured in 2003 ancl 2009.," Our new observation demonstrates that within two months after the cessation of the recent accretion outburst, the thermal flux and neutron star temperature are elevated above the quiescent base level measured in 2003 and 2009." + In analogy with that seen for quasi-persistent LAINBs. we attribute this o heating of the neutron star crust due to its bright 2010 accretion outburst.," In analogy with that seen for quasi-persistent LMXBs, we attribute this to heating of the neutron star crust due to its bright 2010 accretion outburst." + Hf true. this leads to the clear prediction hat the crust is expected to cool down in the next months. until thermal equilibrium with the core is re-established and he neutron star returns to its quiescent base level.," If true, this leads to the clear prediction that the crust is expected to cool down in the next months, until thermal equilibrium with the core is re-established and the neutron star returns to its quiescent base level." + oobservations carried out within the next vear are thus expected to detect a decrease. in. neutron. star. effective emperature and thermal luminosity. down to the values measured in archival data.," observations carried out within the next year are thus expected to detect a decrease in neutron star effective temperature and thermal luminosity, down to the values measured in archival data." + An alternative explanation that may be invoked. is hat the elevated. emission. level is. caused by residual accretion., An alternative explanation that may be invoked is that the elevated emission level is caused by residual accretion. + Lt has been proposed. that low-level accretion onto the neutron star surface produces thermal radiation (Zampierictal.1995).. and the X-ray spectrum may appear indistinguishable from a passive neutron star. atmosphere (Soriaοἱal.2011).," It has been proposed that low-level accretion onto the neutron star surface produces thermal radiation \citep[][]{zampieri1995}, and the X-ray spectrum may appear indistinguishable from a passive neutron star atmosphere \citep[][]{soria2011}." +. However. there are also indications that ow-level accretion involves a hard spectral component.," However, there are also indications that low-level accretion involves a hard spectral component." + One example is462.. which has been extensively monitored in quiescence withChandra. aand ssince its 1.5-vr lone outburst ended in 2007. (CEridrikssonetal.9010. 2011).," One example is, which has been extensively monitored in quiescence with, and since its 1.5-yr long outburst ended in 2007 \citep[][]{fridriksson2010,fridriksson2011}." +. This source exhibits X-ray Dares. suggestive of (sporadic) low-level aceretion. which are associated with a strong increase in the powerlaw spectral component (Eridrikssonetal.2011).," This source exhibits X-ray flares, suggestive of (sporadic) low-level accretion, which are associated with a strong increase in the powerlaw spectral component \citep[][]{fridriksson2011}." +. In case of JITAS. we detect a purely thermal spectrum. both in the new oobservation anc in archival cata.," In case of J1748, we detect a purely thermal spectrum, both in the new observation and in archival data." + Any possible hard spectral component contributes 154 to the total unabsorbed 0.510 keV Dux., Any possible hard spectral component contributes $\lesssim15\%$ to the total unabsorbed 0.5–10 keV flux. + Furthermore. the observed quicscent thermal emission may vary a factor of ~23 from one quiescent epoch to another. due to the changing amount of fucl anc hvdrogen/helium abuncances present alter cach outburst.," Furthermore, the observed quiescent thermal emission may vary a factor of $\sim2-3$ from one quiescent epoch to another, due to the changing amount of fuel and hydrogen/helium abundances present after each outburst." + This influences the fraction. of heat conducted. from. the crust towards the core and surface. and hence the observec thermal quiescent emission (Brownetal.2002).," This influences the fraction of heat conducted from the crust towards the core and surface, and hence the observed thermal quiescent emission \citep[][]{brown2002}." +. Hf this is the case for J1748. the quiescent thermal Luminosity is expectec to remain at the value inferred from our 2011 oobscrvation. provided that the accretion has switched oll completely.," If this is the case for J1748, the quiescent thermal luminosity is expected to remain at the value inferred from our 2011 observation, provided that the accretion has switched off completely." + We consider crustal heating the most likely explanation for the observed elevated. quiescent emission of the 11 112 X-ray pulsar in Terzan 5. but additional. oobscrvations are required to rule out the above alternative scenarios.," We consider crustal heating the most likely explanation for the observed elevated quiescent emission of the 11 Hz X-ray pulsar in Terzan 5, but additional observations are required to rule out the above alternative scenarios." + For the purpose of calculating crust cooling curves. it is necessary to constrain the end of the accretion outburst.," For the purpose of calculating crust cooling curves, it is necessary to constrain the end of the accretion outburst." + We therefore try to determine the time at which the source intensity cropped below the detection threshold of (implying a 230 keV luminosity of Ly~107eres1)., We therefore try to determine the time at which the source intensity dropped below the detection threshold of (implying a 2–30 keV luminosity of $L_X \sim10^{35}~\lum$ ). + Unfortunately. Sun angle constraints deprived our view of ‘Verzan 5 in 2010 December. so that the end of the outburst was not observed. (see Fig. 132).," Unfortunately, Sun angle constraints deprived our view of Terzan 5 in 2010 December, so that the end of the outburst was not observed (see Fig. \ref{fig:maxi}) )." + The outburst conimenced around 2010. October 10 and the source. was still active on 2010 December 4., The outburst commenced around 2010 October 10 and the source was still active on 2010 December 4. + This indicates a minimum outburst duration of 7 weeks (55 davs)., This indicates a minimum outburst duration of 7 weeks (55 days). + Since no activity was detected when the mmonitoring observations resumed on 2010 December 28. the outburst duration is constrained to be 4. and select the largest correlation in eachtime slice.," To define a signal to noise ratio $SNR_i$ for each slice $i$, we proceed with the normalized Pearson coefficients where $\sigma_i$ denotes the standard deviation of $\rho_i(\delta)$ as a function of $\delta$, and define In order to extract a trajectory in the time frequency domain, we proceed with correlations $\hat{\rho}_i(\delta)>4$, and select the largest correlation in eachtime slice." + This filler extracts at most 52 data points Irom each frame of 52 seconds when 7=, This filter extracts at most 52 data points from each frame of 52 seconds when $\tau=1$ . +ü We derive the average value for the tidal racius à; of galaxies orbiting inside a host halo with radius R and circular velocity V., We derive the average value for the tidal radius $r_{tid}$ of galaxies orbiting inside a host halo with radius $R$ and circular velocity $V$. + Such a radius is that appropriate for a galactic sublalo which survives the tidal stripping of the host halo. a condition which requires the density of the ealactic subhalo within 74; to exceed the deusitv of the host halo iuterior to the poriceutre ry of its orbit.," Such a radius is that appropriate for a galactic subhalo which survives the tidal stripping of the host halo, a condition which requires the density of the galactic subhalo within $r_{tid}$ to exceed the density of the host halo interior to the pericentre $r_p$ of its orbit." + To this cud we adopt the approach of Cduigna ct al. (, To this end we adopt the approach of Ghigna et al. ( +1998) who showed how the condition for survival against tidal stripping trauslates approsimately iuto mi7ry(esi).,"1998) who showed how the condition for survival against tidal stripping translates approximately into $r_{tid}\approx r_p\,(v/V)$." + The above expression for rg; has been tested by the above authors against lieh resolution N-body simulations. and has been proved to agree with the values of 7; measured im the simulatious for all sublialoes except the few on very eccentric orbits (the latter have measured radii larger than expected partly due to the formation of tidal tails).," The above expression for $r_{tid}$ has been tested by the above authors against high resolution N-body simulations, and has been proved to agree with the values of $r_{tid}$ measured in the simulations for all subhaloes except the few on very eccentric orbits (the latter have measured radii larger than expected partly due to the formation of tidal tails)." + We average the above relation for rg; over the distribution of pericentres obtained by Cigna et al. (, We average the above relation for $r_{tid}$ over the distribution of pericentres obtained by Ghigna et al. ( +"1998) from N-body sinmnlatious. which we fit with a modified lognormal expression Gn the variable (7,/R0.08) with logarithmic mica -1.3 and logarithmic variance 0.6).","1998) from N-body simulations, which we fit with a modified lognormal expression (in the variable $(r_p/R-0.08)$ with logarithmic mean -1.3 and logarithmic variance 0.6)." + Performing the average vields for ry; the following expression: where the lower Bunt r4; corresponds to the minima tidal radius that a ealactic subhalo can have without being severely distorted or disrupted., Performing the average yields for $r_{tid}$ the following expression: where the lower limit $r_{cut}$ corresponds to the minimum tidal radius that a galactic subhalo can have without being severely distorted or disrupted. + Following Bullock. Iravstov and Weinberg (2000). we adopt for r.;; the radius of the peak of the ealaxy velocity profile.," Following Bullock, Kravstov and Weinberg (2000), we adopt for $r_{cut}$ the radius of the peak of the galaxy velocity profile." + For à Navarro ct al. (, For a Navarro et al. ( +1997) circular velocity profile rng=2.16ου6 holds. where e is the concentration parameter of the subhlialo. and royy is the radius where the average density of the subhlialo would equal 200p...,"1997) circular velocity profile $r_{cut}=2.16\,r_{200}/c$ holds, where $c$ is the concentration parameter of the subhalo, and $r_{200}$ is the radius where the average density of the subhalo would equal $200\,\rho_c$." + The equations relating e aud rayy) with the circular velocity e of the subhalo are given by the above authors., The equations relating $c$ and $r_{200}$ with the circular velocity $v$ of the subhalo are given by the above authors. + The validity of the above value for (74; has been tested against N-body results by Bullock et al. (, The validity of the above value for $r_{cut}$ has been tested against N-body results by Bullock et al. ( +2000).,2000). + It is easy to recast the relation rg=àRe/V iu terms of deusity.," It is easy to recast the relation $r_{tid}=\alpha\,R\,v/V$ in terms of density." + Substituting the relation e/V.=Gn/M)V?(Rfrg)? one obtains (rgZR)!=a?in/M so that the ratio of the sublalo to the host halo average deusities is eiven by 1/60?.," Substituting the relation $v/V=(m/M)^{1/2}\,(R/r_{tid})^{1/2}$ one obtains $(r_{tid}/R)^3=\alpha^2\,m/M$ so that the ratio of the subhalo to the host halo average densities is given by $1/\alpha^2$." + The average over the distribution of periceutres defining the value of a yields typical values for such a ratio ranging between 3 and 5. depending on the coucentration paraucter e of the sublialoes.," The average over the distribution of pericentres defining the value of $\alpha$ yields typical values for such a ratio ranging between 3 and 5, depending on the concentration parameter $c$ of the subhaloes." + Note that taking the tidal radius rg;=rjc/V. as the limiting radius of the zubhaloes is fully appropriate only for sublialoes mich simaller than their host halo.," Note that taking the tidal radius $r_{tid}=r_p\,v/V$ as the limiting radius of the subhaloes is fully appropriate only for subhaloes much smaller than their host halo." + This is the case for the bulk of the population of satellite ealaxies (the one mainly coutrvibuting to binary ageregations)oOo for which our treatimuenut constitutes a valid approximation., This is the case for the bulk of the population of satellite galaxies (the one mainly contributing to binary aggregations) for which our treatment constitutes a valid approximation. +em level.,cm level. +" Thus. we use (hese measurements as upper limits for the position wander of the centroid of ""ss emission."," Thus, we use these measurements as upper limits for the position wander of the centroid of s emission." +" The position wander of oon time scales of days to weeks was determined from ""dailv position measurements. obtained from z5 hours of data when source elevations were above ez20° al most of the five inner-VLBA antennas."," The position wander of on time scales of days to weeks was determined from “daily” position measurements, obtained from $\approx5$ hours of data when source elevations were above $\approx20^\circ$ at most of the five inner-VLBA antennas." + These positions were plotted in Fig., These positions were plotted in Fig. + 3 of (2004).. after removing the long-term proper motion of oof 6.379 ! along a position angle of 209.60° East of North.," 3 of \citet{Reid-Brun:2004}, after removing the long-term proper motion of of $6.379$ $^{-1}$ along a position angle of $209.60^\circ$ East of North." + We differenced. the EW and NS positions for all pairs of measurements separated by less than 10? min: the nagnitudes of these position differences are shown in Fig. 2..," We differenced the EW and NS positions for all pairs of measurements separated by less than $10^5$ min; the magnitudes of these position differences are shown in Fig. \ref{wander}," + along with weighted averages or (wo time-bins., along with weighted averages for two time-bins. + The average difference in EW position lor measurements separated by zz6000 min (84 davs) is about 400µας., The average difference in EW position for measurements separated by $\approx6000$ min $\approx4$ days) is about $400~\muas$. + This difference is likely not a property of the emission ofA*:: instead it is probably caused. by mus-modeling large-scale propagation delavs (hrough the Earth's atmosphere., This difference is likely not a property of the emission of; instead it is probably caused by mis-modeling large-scale propagation delays through the Earth's atmosphere. + Most of the data used in (his analysis were collected before we started {ο measure atmospheric path-leneths in 2003., Most of the data used in this analysis were collected before we started to measure atmospheric path-lengths in 2003. + Without this calibration. residual zenith length errors of ~5 em are typical for the VLBA correlator model at 7-mm wavelength.," Without this calibration, residual zenith path-length errors of $\sim5$ cm are typical for the VLBA correlator model at 7-mm wavelength." + since propagation delavs can be correlated with large-scale weather patterns. which have characteristic Ue scales of several davs. one expects to see such a residual signature in our position differences.," Since propagation delays can be correlated with large-scale weather patterns, which have characteristic time scales of several days, one expects to see such a residual signature in our position differences." + Limits upon (he variability in the centroid position of high resolution images of Ser A* imply a constraint upon orbiting hot-spot models for Ser A*’s radio variability., Limits upon the variability in the centroid position of high resolution images of Sgr A* imply a constraint upon orbiting hot-spot models for Sgr A*'s radio variability. + This will necessarily be a function of the hot-spot orbit and the hot-spot/clisk flux ratio., This will necessarily be a function of the hot-spot orbit and the hot-spot/disk flux ratio. + In this section we derive the expected. variability for (wo simple hot-spot models., In this section we derive the expected variability for two simple hot-spot models. + The first is an oplically thin Newtonian hot-spot. while the second fully incorporates general relativity ancl the opacity of a disk constimmeted such that it reproduces the observed spectrum of Ser À*.," The first is an optically thin Newtonian hot-spot, while the second fully incorporates general relativity and the opacity of a disk constructed such that it reproduces the observed spectrum of Sgr A*." +and Differentiating Eq. (,and Differentiating Eq. ( +"13) and substituting xBo)+Bo(VB,) mio il give We assume that (1) the accretion disc is infinitesimally thin. (2) the magnetic moment and the spin of the NS are parallel to the z axis. and normal to the disc. i.e... Bg=(0.0.By) and Q=(0.0.0) close to the inner edge of the disc. and (3) the 7 and y axes are along the disc plane. ancl the MIID wave is assumed to propagate in the coz plane. i.e.. (he wave vector ke=(ksin8.0.kcos 8). where @ is the angle between the z axis and &.","13) and substituting $\nabla (\bm {B_{0}} \cdot \bm {B_{s}} ) = (\bm {B_{0}} \cdot +\nabla )\bm {B_{s}} + (\bm {B_{s}} \cdot \nabla )\bm {B_{0}} + \bm +{B_{s}} \times (\nabla \times \bm {B_{0}} ) + \bm {B_{0}} \times +(\nabla \times \bm {B_{s}} )$ into it give We assume that (1) the accretion disc is infinitesimally thin, (2) the magnetic moment and the spin of the NS are parallel to the $z$ axis, and normal to the disc, i.e., ${ {\bm B}}_{ {\bm 0}} = +(0,0,B_0 )$ and ${ {\bm \Omega}}=(0,0,\Omega )$ close to the inner edge of the disc, and (3) the $x$ and $y$ axes are along the disc plane, and the MHD wave is assumed to propagate in the $xoz$ plane, i.e., the wave vector ${ {\bm k}}= (k\sin \theta ,0,k\cos \theta )$ , where $\theta $ is the angle between the $z$ axis and $ {\bm k}$." + After carrving out. Fourier transformation lor Eqs. (, After carrying out Fourier transformation for Eqs. ( +"14)-(16) we get the following dispersion equations. where / is imaginary unit. V4 (=V Basin). e, (= vom). wy, (=VOM η). and w are Alfvénn velocity. acoustic velocity. IXeplerian angular velocity. and angular velocity of the perturbation at ry respectively. and. oy). c4. 0; avethe three components of the perturbed quantity of the speed.","14)-(16) we get the following dispersion equations, where $i$ is imaginary unit, $V_A$ $= \sqrt {B_0 ^2 / \mu \rho _0 +}$ ), $c_s$ $= \sqrt {\gamma P_0 / \rho_0}$ ), $\omega _k$ $= \sqrt +{GM / r_0 ^3}$ ), and $\omega$ are Alfvénn velocity, acoustic velocity, Keplerian angular velocity, and angular velocity of the perturbation at $r_0$ respectively, and $v_{sx}$, $v_{sy}$, $v_{sz}$ arethe three components of the perturbed quantity of the speed." + Equations (17)-(19) show that there exist three resonance ALLID modes., Equations (17)-(19) show that there exist three resonance MHD modes. + At the magnetospheric radius rj. the magnetic energy density is equal to the total kinetic energy density. ie. B?/S8a=pV1/2ep¥2/2 (Davidson Ostriker 1973:Ghosh et al.," At the magnetospheric radius $r_0$, the magnetic energy density is equal to the total kinetic energy density, i.e. $ B^2 / 8\pi = \rho +V_A^2/2 \simeq \rho V_k^2/2$ (Davidson Ostriker 1973;Ghosh et al." + 1977)., 1977). + Since the characteristic wavelength is in the same order with the magnetospheric radius (Rezania Samson 2005). we then have &KV4~ΑνVyfry=uw. or KV=e.," Since the characteristic wavelength is in the same order with the magnetospheric radius (Rezania Samson 2005), we then have $kV_A \sim kV_K \sim V_K/r_0 = \omega_k $, or $kV_A = \eta +\omega _k$." + Because the thermal pressure of the plasma might be comparable with the magnetic pressure (c;~ V4) just inside the magnetosphere (Miller et al., Because the thermal pressure of the plasma might be comparable with the magnetic pressure $c_s \sim V_A$ ) just inside the magnetosphere (Miller et al. +" 1993). we also suppose fre,=Aw. Ile"," 1998), we also suppose $kc_s = +\lambda \omega _k$." +re both η ancl A are taken (o be constant for certain sources., Here both $\eta$ and $\lambda$ are taken to be constant for certain sources. + Substitute these relations into Eqs. (, Substitute these relations into Eqs. ( +17)-(19) we can get the resonant modes.,17)-(19) we can get the resonant modes. + Specificallywhen 9=0. from Eqs. (," Specificallywhen $\theta = 0$, from Eqs. (" +"17) (18) we can gel vy, =cv.∙ Le. releHer belhw—HerUu3holsisand .","17) and (18) we can get $v_{sx} = \pm +iv_{sy}$ , i.e. $v_{sx} e^{i{\bm k} \cdot {\bm r} - i\omega t} = +v_{sy} e^{i{\bm k} \cdot {\bm r} - i\omega t\pm i{\textstyle{{\pi} +\over {2}}}}$ ." + substituting this relation into the Eqs. (, Substituting this relation into the Eqs. ( +17)-(19) can give,17)-(19) can give +magnetic flux. produced in the torus.,magnetic flux produced in the torus. + This lorque received from the black hole compensates [or angular momentum losses in magnetic wiudsa d raclialou. which arrests the inflow aud enables a state of suspended accretion around a rapicly Spinuing black hole.," This torque received from the black hole compensates for angular momentum losses in magnetic winds and radiation, which arrests the inflow and enables a state of suspended accretion around a rapidly spinning black hole." + The suspended accretion state is expected to be stable on average (vanPue&Ost‘iker2000)., The suspended accretion state is expected to be stable on average \citep{mvp00c}. +. There powerful shear between the inier and the «titer faces of the torus is. to leacdiug order. dictated by Ixeplerian. motion.," There powerful shear between the inner and the outer faces of the torus is, to leading order, dictated by Keplerian motion." + Some deviaüOl away rom Weplerian motion is expected due to the presence of competing torques., Some deviation away from Keplerian motion is expected due to the presence of competing torques. + They eds to brit[n] ie. two [aces in state of super- aud motion. with positive radial preSS' which p‘omotes a slender shape.," They tends to bring the two faces in state of super- and sub-Keplerian motion, with positive radial pressure which promotes a slender shape." + The interlace separating the two [aces is expected to be 1istable. wuch favors turbulent imixiug into a state of Που specific euergy across the torus.," The interface separating the two faces is expected to be unstable, which favors turbulent mixing into a state of uniform specific energy across the torus." +" Mixiug euliances differential rotatiol. as may be illustrated in the Newtonian limit. which gives rise to he angular veocity O(r)zzOn,(1—(rαλα): ""uasa duction of raclit sr fora torus of major radius e."," Mixing enhances differential rotation, as may be illustrated in the Newtonian limit, which gives rise to the angular velocity $\Omega(r)\approx\Omega_K(1-(r-a)/a)^{1/2}$ as a function of radius $r$ for a torus of major radius $a$." + Compression into a mo'e slender shape tends to 'educe clillerential rotation., Compression into a more slender shape tends to reduce differential rotation. + The net result s10uld. be tllat he characteristically Ixeplerian decrease of angular velocity with radius is approximately preserved., The net result should be that the characteristically Keplerian decrease of angular velocity with radius is approximately preserved. +" The inner aid other faces will have. 'espectively. augular velocities ο2Op,(1d:36/La). whe'e b denotes their radial separation."," The inner and other faces will have, respectively, angular velocities $\Omega_\pm\approx \Omega_K\left(1\pm 3b/4a\right),$ where $b$ denotes their radial separation." + The salue trend should hold in the Ίου metric., The same trend should hold in the Kerr metric. + Craviatioual radiation from a torus surrouudiug a black hole teuds to domiuate radio waves of he same frequency., Gravitational radiation from a torus surrounding a black hole tends to dominate radio waves of the same frequency. + This is gene‘ally due to tle COLipact size in the presence of gravitationally Weruk magnetic fields., This is generally due to the compact size in the presence of gravitationally weak magnetic fields. +" Consider a orus with elipticiy €, a magnetic moment sop and mass im in rotation about its center of mass."," Consider a torus with ellipticity $\epsilon$, a magnetic moment $\mu_T$ and mass $m$ in rotation about its center of mass." +" Its ¢1adrupolar motilents iu magnetic moment aud ass are. respectively. eji aud em. which produce lujnosities (acapted [rom ShapiroaudTeukolsky (1983))): LomSH(OpM)μη:f-)-eydD> aud Law72(32/5(Op.1)""(nM)ein ⋅geometrical u"" losityits."," Its quadrupolar moments in magnetic moment and mass are, respectively, $\epsilon \mu$ and $\epsilon m$, which produce luminosities (adapted from \cite{sha83}) ): ${\cal L}_{em}\approx\pi^{-1}(\Omega_T M)^4(\mu_T/M^2)^2\epsilon^2$ and ${\cal L}_{gw}\approx ({32}/{5})(\Omega_T M)^{10/3} (m/M)^2\epsilon^2$ in geometrical units." +" These ellissious may be compared with. fespecively. he ltniiu radio emission ~ Qi~/pi2g from ali orthogoual pulsar zu in gravitajonal-wave enisslons (32/5)(0,,5,M1from ne Urol star-leltron star binaries wih angular velociy Ou aud chi‘p mass M=(AlpAlo)?/(M+Aly) (for circular orbits)."," These emissions may be compared with, respectively, the luminosity in radio emission $\sim \Omega^4_p\mu_p^2/\pi$ from an orthogonal pulsar and in gravitational-wave emissions $\sim ({32}/{5})(\Omega_{orb} {\cal M})^{10/3}$ from neutron star-neutron star binaries with angular velocity $\Omega_{orb}$ and chirp mass ${\cal M}=(M_1M_2)^{3/5} /(M_1+M_2)^{1/5}$ (for circular orbits)." +" The ratio of radio-to-gravitationale wave emissions cau be evaltuatec als e.g..ο) when Ep/M~10"" for the relative energy in the magnetic field and M1»p102."," The ratio of radio-to-gravitational wave emissions can be evaluated as e.g., when $E_B/M\sim 10^{-6}$ for the relative energy in the magnetic field and $M/m\le 10^2$." + The fluence Fey) in gravitational radiatiou may be appreciable. provided the torus is long-lived iu a state of suspended accretion for the duraction of the burst.," The fluence $F_{GW}$ in gravitational radiation may be appreciable, provided the torus is long-lived in a state of suspended accretion for the duraction of the burst." + The suspended accretion state is described by balauce of torque aud euerey:, The suspended accretion state is described by balance of torque and energy: +trend that the diminiug of light appears to end first at louger wavelengtlis.,trend that the dimming of light appears to end first at longer wavelengths. + However. this needs to," However, this needs to" +These periocdicities are dillerent from the ones detected in their pulse arrival times except for the 1260235. day periodicity of PSR 1458. which is close to the pulse arrival time periodicity of 1186437. days.,"These periodicities are different from the ones detected in their pulse arrival times except for the $1260^{+386}_{-240}$ day periodicity of PSR $-$ 1458, which is close to the pulse arrival time periodicity of $\pm$ 7 days." + In general. the pulse arrival time and. pulse rate recurrence periodicities are expected to be independent though it is possible for them. to be the same., In general the pulse arrival time and pulse rate recurrence periodicities are expected to be independent though it is possible for them to be the same. + In order to further gauge the reliability of our resultsfor PSRs 1458 and 30. we created 1000 rancom sequences by assigning the J1754.—measured: pulse detection rates to randomly. selected ALJDs and calculated the number of times a peak of the same or higher significance appeared in the 1000 trials.," In order to further gauge the reliability of our results for PSRs $-$ 1458 and $-$ 30, we created 1000 random sequences by assigning the measured pulse detection rates to randomly selected MJDs and calculated the number of times a peak of the same or higher significance appeared in the 1000 trials." + This number was six and ten for, This number was six and ten for +"In recent years, multi-waveleneth selection techniques have been successful in identifying different populations of star-forming galaxies at high redshifts, :—6 [see Ellis(2008) for a recent review].","In recent years, multi-wavelength selection techniques have been successful in identifying different populations of star-forming galaxies at high redshifts, $z \gtrsim 6$ [see \citet{ellis08} for a recent review]." +" An important class of such galaxies consists of the Lyman-o emitters (hereafter LAEs), objects identified through their excess emission in narrow-band images centred on the redshifted Lyman-c wavelength (c.g. Huetal.1998;Rhoads2000;Taniguchi 2005))."," An important class of such galaxies consists of the $\alpha$ emitters (hereafter LAEs), objects identified through their excess emission in narrow-band images centred on the redshifted $\alpha$ wavelength (e.g. \citealp{hu98,rhoads00,taniguchi05}) )." +" Follow-up Lyman-a spectroscopy has yielded accurate redshifts for a significant fraction of the LAE population (c.g. Taniguchietal.2005:Kashikawa 2006)), unlike most other high-- star-forming galaxies (c.9. Lyman-break systems, sub-mm galaxies), which typically only have photometric redshifts."," Follow-up $\alpha$ spectroscopy has yielded accurate redshifts for a significant fraction of the LAE population (e.g. \citealp{taniguchi05,kashikawa06}) ), unlike most other $z$ star-forming galaxies (e.g. Lyman-break systems, sub-mm galaxies), which typically only have photometric redshifts." + In fact. an LAE at 2=6.96 (yeetal.2006) has the highest confirmed spectroscopic redshift of all presently-known galaxies.," In fact, an LAE at $z = 6.96$ \citep{iye06} has the highest confirmed spectroscopic redshift of all presently-known galaxies." +" LAEs constitute a significant fraction of the star-forming galaxy population at 2~6, sufficient to reionize the Universe at earlier epochs (e.g. Fanetal. 2006))."," LAEs constitute a significant fraction of the star-forming galaxy population at $z \sim 6$, sufficient to reionize the Universe at earlier epochs (e.g. \citealp{fan06}) )." + The number density of LAEs and the typical shape of the Lyman-o emission line therefore provide important probes of physical conditions in the Universe around the epoch of reionization (e.g. Haiman&Spaans1999;2002;Kashikawaetal. 2006)).," The number density of LAEs and the typical shape of the $\alpha$ emission line therefore provide important probes of physical conditions in the Universe around the epoch of reionization (e.g. \citealp{haiman99,haiman02,kashikawa06}) )." +" Equally important, the steep drop in the space density of quasars at >=6 implies that they would be unable to produce the ultraviolet (UV) background radiation required to reionize the Universe (e.g. Fanetal. 2001))."," Equally important, the steep drop in the space density of quasars at $z \gtrsim 6$ implies that they would be unable to produce the ultraviolet (UV) background radiation required to reionize the Universe (e.g. \citealp{fan01}) )." + Star-forming galaxies like the LAEs are thus most likely to have been responsible for reionization (c.g. 2004))., Star-forming galaxies like the LAEs are thus most likely to have been responsible for reionization (e.g. \citealp{yan04}) ). +" Understanding the factors that influence the star-formation activity in these systems (e.g. the molecular gas content, the star formation efficiency, etc) is hence of much importance."," Understanding the factors that influence the star-formation activity in these systems (e.g. the molecular gas content, the star formation efficiency, etc) is hence of much importance." + Recent studies of individual LAEs at 4.2<26.6 have shown that these galaxies have a, Recent studies of individual LAEs at $4.2 < z < 6.6$ have shown that these galaxies have a +Emission from pair cascades initiated by photopion production of ultra-high-energy protons in GRB outflows can yield spectra with roughly equal power over a broad energy range.,Emission from pair cascades initiated by photopion production of ultra-high-energy protons in GRB outflows can yield spectra with roughly equal power over a broad energy range. + While GeV-TeV signatures due to hadronic processes have been discussed previously, While GeV-TeV signatures due to hadronic processes have been discussed previously +Websteretal.(1995) compared the £dy colours of Parkes UallJansky Flat-Spectrum Sample (PLIES) (Drinkwateretal.L997) quasars to those of Large Bright Quasar Survey (LBQS) (Llewett.Foltz.&Challee1995) quasars.,\citet{1995Natur.375..469W} compared the $B - K$ colours of Parkes Half-Jansky Flat-Spectrum Sample (PHJFS) \citep{1997MNRAS.284...85D} quasars to those of Large Bright Quasar Survey (LBQS) \citep{1995AJ....109.1498H} quasars. + The PLLJES. quasars were found. to have a significantly ποσο ane redder distribution. of Dἐν colours than LBQS quasars., The PHJFS quasars were found to have a significantly broader and redder distribution of $B - K$ colours than LBQS quasars. + This disparity should not exist according to 1e unified moclel ofActive Galactic Nuclei (AGN) ealaxies (Antonueci1903)..., This disparity should not exist according to the unified model ofActive Galactic Nuclei (AGN) galaxies \citep{1993ARA&A..31..473A}. + Websteretal.(1995). concluded. that rere exists a population of red quasars that at that time aac not been discovered., \citet{1995Natur.375..469W} concluded that there exists a population of red quasars that at that time had not been discovered. + Websteretal.(1995) investigated js possibility using a simple model of LBQS selection Iects. and demonstrated that the LDOS was biased against 1e detection. of red. quasars.," \citet{1995Natur.375..469W} investigated this possibility using a simple model of LBQS selection effects, and demonstrated that the LBQS was biased against the detection of red quasars." + The bias was caused. by the due magnitude limit of the LBOS. because in a blue-lux imited survey. red quasars need to be intrinsically brighter iun blue quasars to be detected. (Websterctal.1995).," The bias was caused by the blue magnitude limit of the LBQS, because in a blue-flux limited survey, red quasars need to be intrinsically brighter than blue quasars to be detected \citep{1995Natur.375..469W}." + The bias causes fewer red quasar detections. because the uminosity function of quasars is steep at the bright. end Dovle.Jones.&Shanks1991)..," The bias causes fewer red quasar detections, because the luminosity function of quasars is steep at the bright end \citep{1991MNRAS.251..482B}." + Other surveys with a blue-lux limit are similarly. biassed. independent. of the quasar selection method.," Other surveys with a blue-flux limit are similarly biassed, independent of the quasar selection method." + Using their simple model. estimated that the red. quasars missed by. previous surveys constitute as much as of all quasars.," Using their simple model, estimated that the red quasars missed by previous surveys constitute as much as of all quasars." + Rec quasars are defined by a colour criterion., Red quasars are defined by a colour criterion. + The best colour to use for this criterion. is that made from filters separated by the largest wavelength range possible. because anomalously redder quasar colours stand out more.," The best colour to use for this criterion is that made from filters separated by the largest wavelength range possible, because anomalously redder quasar colours stand out more." + For the datasets used in this paper. the colour covering the largest wavelength range is (dy.," For the datasets used in this paper, the colour covering the largest wavelength range is $U - K$." + We used the next largest. colour b;—dv. as far more objects. ~ 3 times as many. have a measured b;/—A colour compared to the number with a measurecl £CI colour.," We used the next largest colour $b_J - K$, as far more objects, $\sim$ 3 times as many, have a measured $b_J - K$ colour compared to the number with a measured $U - K$ colour." + The additional advantage of having used. b;dv ds that it is possible to compare the colours of the quasar sample presented here with the DΑν colour used in Clikmanetal.(2004) and the g9A colour of Hopkinsetal. (2004).., The additional advantage of having used $b_J - K$ is that it is possible to compare the colours of the quasar sample presented here with the $B - K$ colour used in \citet{2004ApJ...607...60G} and the $g - K$ colour of \citet{2004AJ....128.1112H}. +" Throughout this paper we define a red quasar by the criterion. b—Azx 3.5 (Cireggetal.2002). because when describing the optical-infrared region of quasar spectra with a power-law. S(A)-xA"". this colour corresponds to an ellective optical-infrared spectral index of azx1."," Throughout this paper we define a red quasar by the criterion, $b_J - K \geq$ 3.5 \citep{2002ApJ...564..133G}, because when describing the optical-infrared region of quasar spectra with a power-law, $S(\lambda) \propto \lambda^{\alpha}$, this colour corresponds to an effective optical-infrared spectral index of $\alpha \geq 1$." + The existence. of red. quasars. other than in. radio-selected quasar samples. over most of the b;fy colour range of the PLLJES quasars. satisfving the above criterion. has already been demonstrated. (Hopkinsetal.2004: 2004).. ," The existence of red quasars, other than in radio-selected quasar samples, over most of the $b_J - K$ colour range of the PHJFS quasars, satisfying the above criterion, has already been demonstrated \citep{2004AJ....128.1112H, 2003AJ....126.1131R, 2004ApJ...607...60G}. ." +Phe Clikmanetal.(2004) τος quasar sample contains the reddest quasars found. with D N colours of 5xBpΑν S.," The \citet{2004ApJ...607...60G} red quasar sample contains the reddest quasars found, with $B - K$ colours of $5 \leq B - K \leq 8$ ." +seen in real data assuming the same density of background SOULCOS.,seen in real data assuming the same density of background sources. + When carrving out a direct reconstruction. the first issue to be aclelressecl is tha of the appropriate choice of bin size over which to average the Hexion signal.," When carrying out a direct reconstruction, the first issue to be addressed is that of the appropriate choice of bin size over which to average the flexion signal." + A balance must be struck between using a fine grid. in order to resolve structures on both small anc large scales. and having a statistically significant. number of background: sources within each bin.," A balance must be struck between using a fine grid, in order to resolve structures on both small and large scales, and having a statistically significant number of background sources within each bin." +" In Figure Y.. we srow the fraction of pixels containing at least μα sources as à function of bin width for Ny,=1 and 3."," In Figure \ref{fg:frac_covered}, we show the fraction of pixels containing at least $N_{\rm min}$ sources as a function of bin width for $N_{\rm min}=1$ and $3$." + To alleviate additional pixel noise in our reconstruction resulting [from incomplete coverage. we choose à bin size of O48aremin per pixel. resulting. in à reconstructed convergence map of 16 x 16 pixels.," To alleviate additional pixel noise in our reconstruction resulting from incomplete coverage, we choose a bin size of $0.43\,$ arcmin per pixel, resulting in a reconstructed convergence map of 16 x 16 pixels." + In order |to reduce edge elfects when carrving out the Fourier transforms. we acd zero-padcding on alb four sides of the image such that the padded image consisted of 128 x 128 pixels.," In order to reduce edge effects when carrying out the Fourier transforms, we add zero-padding on all four sides of the image such that the padded image consisted of 128 x 128 pixels." + To estimate the noise in our reconstruction. the Uexion vectors were rotated by a random angle. and the reconstruction repeated.," To estimate the noise in our reconstruction, the flexion vectors were rotated by a random angle, and the reconstruction repeated." + This procedure was carried out. 1000 times. and the error at each pixel location estimated by the standard deviation o; of the randomised reconstructions.," This procedure was carried out 1000 times, and the error at each pixel location estimated by the standard deviation $\sigma_i$ of the randomised reconstructions." + Finally. the. reconstructed. convergence must be appropriately rescaled.," Finally, the reconstructed convergence must be appropriately rescaled." + As our data covers only the central th! Mpe of the eliister. we do not expect the convergence to &o to zero at the edge of the cluster.," As our data covers only the central $h^{-1}\,$ Mpc of the cluster, we do not expect the convergence to go to zero at the edge of the cluster." + We therefore set the integration constan Bo such that the weighted mean of the reconstructed convergence. defined as: is equal to the mean convergence in the input. map. found to be &=0.13.," We therefore set the integration constant $\kappa_0$ such that the weighted mean of the reconstructed convergence, defined as: is equal to the mean convergence in the input map, found to be $\overline{\kappa}=0.13$." + Figure 8. shows the convergence reconstructed. in this manner. as well as à comparison of the radial profile of the reconstruction with that of the input convergence map.," Figure \ref{fg:directrec} shows the convergence reconstructed in this manner, as well as a comparison of the radial profile of the reconstruction with that of the input convergence map." + Note that the error bars represent the 1-0 errors on hy., Note that the error bars represent the $1$ $\sigma$ errors on $\kappa-\kappa_0$. + Alodelling the convergence as a power law of the form HoxorC7. we find that the convergence reconstruction shows a shallow slope (a= 0.19). while the input map shows a much. steeper slope (a~ 0.66).," Modelling the convergence as a power law of the form $\kappa \propto r^{-\alpha}$, we find that the convergence reconstruction shows a shallow slope $\alpha = 0.19$ ), while the input map shows a much steeper slope $\alpha \sim 0.66$ )." + We note also that the input convergence appears to follow a broken power law. rather than having a constant power law index. whilst the reconstruction does not appear to follow this behaviour.," We note also that the input convergence appears to follow a broken power law, rather than having a constant power law index, whilst the reconstruction does not appear to follow this behaviour." + The underestimation of the central slope is not entirely unexpected., The underestimation of the central slope is not entirely unexpected. + The reason is that one expects a decrease in the number of background. sources with [lexion. measurements in the central regions of the cluster. resulting from sources being lensed. away from the centre. and. from. blending of background images with those of cluster members.," The reason is that one expects a decrease in the number of background sources with flexion measurements in the central regions of the cluster, resulting from sources being lensed away from the centre, and from blending of background images with those of cluster members." + Figure 9 shows the mean background. source count evaluated in annuli as a function of distance from the cluster centre. demonstrating a significant unclerclensity in the central 1.5 arcmünutes as compared. with the periphery.," Figure \ref{fg:source_count} shows the mean background source count evaluated in annuli as a function of distance from the cluster centre, demonstrating a significant underdensity in the central 1.5 arcminutes as compared with the periphery." + “Phe lack of sources in the central region limits the ability of this method to constrain the potential in that regime., The lack of sources in the central region limits the ability of this method to constrain the potential in that regime. + Note that duc to our choice of normalisation constant.," Note that due to our choice of normalisation constant," +The analytic description of the vacuum spacetime surrounding a rapidly rotating neutron star is stil| an open problem.,The analytic description of the vacuum spacetime surrounding a rapidly rotating neutron star is still an open problem. + The analytic structure of the spacetime outside a slowly rotating star. and its relation to the Kerr metric. has been well understood since the seminal works of Hartle (1968) and Hartle Thorne (1969).," The analytic structure of the spacetime outside a slowly rotating star, and its relation to the Kerr metric, has been well understood since the seminal works of Hartle (1968) and Hartle Thorne (1969)." +" On the other hand. numerical solutions o ""the Einstein equations for stars rotating up to the mass-shedding limit are now routinely obtained with a number of different methocs. such as the Komatsu. Eriguchi and Hachisu (1989) method (see Stergioulas 2003. for an extensive comparison of the different exising numerical methods)."," On the other hand, numerical solutions of the Einstein equations for stars rotating up to the mass-shedding limit are now routinely obtained with a number of different methods, such as the Komatsu, Eriguchi and Hachisu (1989) method (see Stergioulas 2003, for an extensive comparison of the different existing numerical methods)." +" These numerical solutions are indeed useful for modeling astrophysica systems. for studying linear perturbations of rapidly rotatinn relativistic stars and as initial data for dynamical evolutions of spacetimes in numerical relativity (see e.g. Stergioulas Friedman 1998. Stergioulas. Kluzniak Bulik 1999, Sergioulas Fon 2001)."," These numerical solutions are indeed useful for modelling astrophysical systems, for studying linear perturbations of rapidly rotating relativistic stars and as initial data for dynamical evolutions of spacetimes in numerical relativity (see e.g. Stergioulas Friedman 1998, Stergioulas, Kluzniak Bulik 1999, Stergioulas Font 2001)." + Despite the availability of numerical soluions. a consisten analytic. representation of the vacuum metric outside a rapidly rotating neutron star is desirable for several reasons.," Despite the availability of numerical solutions, a consistent analytic representation of the vacuum metric outside a rapidly rotating neutron star is desirable for several reasons." + In the firs place. having an analytic. form for the metric simplifies the computation of the properties of the spacetime.," In the first place, having an analytic form for the metric simplifies the computation of the properties of the spacetime." + For example. if an accurate analytic solution were available. geodesics in the neutron star exterior could be studied analytically. and one could tind closed-form expressions for he radii and frequencies of the innermost stable circular orbits WSCOs).," For example, if an accurate analytic solution were available, geodesics in the neutron star exterior could be studied analytically, and one could find closed-form expressions for the radii and frequencies of the innermost stable circular orbits (ISCOs)." + In turn. this would simplify the calculation of properies of accretion disks. of epicyclic frequencies. of accretion lumiysities. and so on.," In turn, this would simplify the calculation of properties of accretion disks, of epicyclic frequencies, of accretion luminosities, and so on." +" Furthermore. having. an analytic. solution. could prove usefu. to the study of properties of the spacetime. such as gravitational"" wave emission."," Furthermore, having an analytic solution could prove useful to the study of properties of the spacetime, such as gravitational wave emission." +s One- of. the unsolved problems .in gravitational-wave theory is the study of he quasinormal modes of rapidly rotating neutron stars., One of the unsolved problems in gravitational-wave theory is the study of the quasinormal modes of rapidly rotating neutron stars. + These can be computed either in the frequency domain. :is an eigenvalue problem. or in the time domain. evolving numericaly the dinearized or full) Einstein equations anc then computing the outgoing radiation.," These can be computed either in the frequency domain, as an eigenvalue problem, or in the time domain, evolving numerically the (linearized or full) Einstein equations and then computing the outgoing radiation." +" The major technical issue in this problem is related to the difficulty of imposing outgoing-wave boundary conditions at infinity. since a rapidly rotating. neutron star Spacetime Is expected to deviate significantly from Petrov type D. Having in hand an accurate analytic metric for the exterior spacetime one could envisage the possibility of computing the Wey! sealars in closed form. looking for neutron star models which are. in some suitably defined sense. “close to Petrov type D"" (Baker"," The major technical issue in this problem is related to the difficulty of imposing outgoing-wave boundary conditions at infinity, since a rapidly rotating neutron star spacetime is expected to deviate significantly from Petrov type D. Having in hand an accurate analytic metric for the exterior spacetime one could envisage the possibility of computing the Weyl scalars in closed form, looking for neutron star models which are, in some suitably defined sense, “close to Petrov type D” (Baker" +The general scenario we envision merely assumes that mass elements that predominantly produce the lighter r-process nuclei are ejected before the mass elements that produce the heavier r-process nuclides.,The general scenario we envision merely assumes that mass elements that predominantly produce the lighter r-process nuclei are ejected before the mass elements that produce the heavier r-process nuclides. + Any setting within a core-collapse supernova that satisfies {his condition would allow for a truncated r-process., Any setting within a core-collapse supernova that satisfies this condition would allow for a truncated r-process. + To evaluate (his idea quantitatively. however. we apply the neutrino-driven wind model for the r-process.," To evaluate this idea quantitatively, however, we apply the neutrino-driven wind model for the r-process." + Although. as noted above. this model is not without issues: it is plausible ancl well discussed in (he literature aud makes a eood selling lor discussing the iruncated r-process.," Although, as noted above, this model is not without issues; it is plausible and well discussed in the literature and makes a good setting for discussing the truncated r-process." + In the neutrino wind scenario of the r-process. neutrinos from (he nascent neutron star heat material at the neutron star surface and drive it away in a wind.," In the neutrino wind scenario of the r-process, neutrinos from the nascent neutron star heat material at the neutron star surface and drive it away in a wind." + As a wind element leaves (he neutron star surface. it expands and cools.," As a wind element leaves the neutron star surface, it expands and cools." + Nucleons in (he wind element assemble into heavier nuclei. which serve as (he seeds [or the subsequent r-process 1994).," Nucleons in the wind element assemble into heavier nuclei, which serve as the seeds for the subsequent r-process \citep{woosley94}." +. Importantly. the abundance pattern for the seed nuclei is well-described by a quasi- which. for the entropy anc neutron richness characteristic of neutrino-driven wind envirnoments. has a relatively sharp abundance peak near A100.," Importantly, the abundance pattern for the seed nuclei is well-described by a quasi-equilibrium, which, for the entropy and neutron richness characteristic of neutrino-driven wind envirnoments, has a relatively sharp abundance peak near $\sim$ 100." + As the temperature in the wind element falls further. charged-particle reactions freeze out.," As the temperature in the wind element falls further, charged-particle reactions freeze out." + If a hieh neutron density accompanies (he seed nuclei. neutron captures ancl beta decays successively promote the seed nuclei to higher mass.," If a high neutron density accompanies the seed nuclei, neutron captures and beta decays successively promote the seed nuclei to higher mass." + This upwards flow in mass slows at the neutron closed shells at 82 and then 126 neutrons. and thereby produces the r-process abundance peaks al A~130 and 195 u (Wooslevetal.1994;Wanajo2002).," This upwards flow in mass slows at the neutron closed shells at 82 and then 126 neutrons, and thereby produces the r-process abundance peaks at $\sim$ 130 and 195 u \citep{woosley94, wanajo02}." +. For what follows. it is essential to note that the neutrino-«lriven wind leaves the star over a period of several seconds.," For what follows, it is essential to note that the neutrino-driven wind leaves the star over a period of several seconds." + In the standard scenario (Meveretal.1992:Woosleyοἱ1994:Takahashi.Witti.&Janka 1994).. the wind elements that leave the star early have lower entropy and a lower degree of neutron richness.," In the standard scenario \citep{meyer92, woosley94, takahashi94}, the wind elements that leave the star early have lower entropy and a lower degree of neutron richness." + As the neutron star evolves. the wind elements that leave later have higher entropy and a greater degree of neutron richness.," As the neutron star evolves, the wind elements that leave later have higher entropy and a greater degree of neutron richness." + It is in (hese latter wind elements that (he heaviest nuclei are made., It is in these latter wind elements that the heaviest nuclei are made. + The lower-entropy. less neutron rich wind elements (hat leave the neutron star earlier produce ligher nuclei.," The lower-entropy, less neutron rich wind elements that leave the neutron star earlier produce ligher nuclei." + If the neutron star collapses to a black hole alter the lower-entropy elements leave the star but before the hieher-entropy do. then the r-process has been truncated. and the abundance pattern will be dominated by lishter r-process nuclei.," If the neutron star collapses to a black hole after the lower-entropy elements leave the star but before the higher-entropy do, then the r-process has been truncated, and the abundance pattern will be dominated by lighter r-process nuclei." + Current supernova models do not naturally produce the parameters required to produce a successful r-process. most notably the entropy. but. also the requisite neutron richness.," Current supernova models do not naturally produce the parameters required to produce a successful r-process, most notably the entropy, but also the requisite neutron richness." + While it may in [act be the case that core-collapse supernovae are not (he site of the r-process. ii may also be (hat current. supernova models are not sufficiently. detailed to provide an," While it may in fact be the case that core-collapse supernovae are not the site of the r-process, it may also be that current supernova models are not sufficiently detailed to provide an" +Hot subdwarf B thereafter: sdB) stars are horizontal branch stars that consist of a helium rich core surrounded by a thin hydrogen envelope.,Hot subdwarf B (hereafter: sdB) stars are horizontal branch stars that consist of a helium rich core surrounded by a thin hydrogen envelope. + The helium core sustains nuclear burning of helium into carbon and oxygen. though in some cases the core may have exhausted helium at the center.," The helium core sustains nuclear burning of helium into carbon and oxygen, though in some cases the core may have exhausted helium at the center." + Their progenitors are main sequence stars that are likely to have undergone the core helium flash C17.ss2A7. Y., Their progenitors are main sequence stars that are likely to have undergone the core helium flash $M \lesssim 2 M_{\odot}$ ). + Once hydrogen in the core is exhausted. their progenitors evolve towards the tip of the red giant branch. where they lose all but a tiny fraction of hydrogen during (or prior to) the helium flash.," Once hydrogen in the core is exhausted, their progenitors evolve towards the tip of the red giant branch, where they lose all but a tiny fraction of hydrogen during (or prior to) the helium flash." + Those stars that retain less than about 17. of hydrogen lie on the high Ziaur end of the horizontal branch., Those stars that retain less than about $M_{\odot}$ of hydrogen lie on the high $T_{\rm eff}$ end of the horizontal branch. + This small fraction of hydrogen has an influence on the future evolution of sdB stars., This small fraction of hydrogen has an influence on the future evolution of sdB stars. + As hydrogen shell burning cannot be sustained by the thin hydrogen envelope. the objects will," As hydrogen shell burning cannot be sustained by the thin hydrogen envelope, the objects will" +Phe method of Schave to characterise [ine-widths is based on Voigt profile fitting of absorption lines (Webb 1987. Carswell 981).,"The method of Schaye to characterise line-widths is based on Voigt profile fitting of absorption lines (Webb 1987, Carswell 1987)." + Phe rationale behind fitting absorption lines with a Voigt profile is partly. historical. and stems from earlier theoretical models in which the forest was produced by a set οἱ.," The rationale behind fitting absorption lines with a Voigt profile is partly historical, and stems from earlier theoretical models in which the forest was produced by a set of." + Phe line-width of these absorbers was assumed to be set by thermal ancl broadening. which would produce a Voig profile. and line blending was responsible for the lines with large deviations from the Voigt profile.," The line-width of these absorbers was assumed to be set by thermal and broadening, which would produce a Voigt profile, and line blending was responsible for the lines with large deviations from the Voigt profile." + In. the new paradigm of the forest absorption in the general LGA is responsible for lines. and there is no a priori reason to expect lines to have the Voigt shape.," In the new paradigm of the forest absorption in the general IGM is responsible for lines, and there is no a priori reason to expect lines to have the Voigt shape." + In this paper we discuss a dillerent method of characterising line-widths. based on discrete wavelets (see c.g. Press et al.," In this paper we discuss a different method of characterising line-widths, based on discrete wavelets (see e.g. Press et al." + 1992 [for an introduction ancl further references)., 1992 for an introduction and further references). + Wavelets provide. an orthogonal basis for a unique decomposition of a signal (the spectrum) in ternis of localised functions with a finite bancwidth., Wavelets provide an orthogonal basis for a unique decomposition of a signal (the spectrum) in terms of localised functions with a finite bandwidth. + Thus they are à compromise between characterising a signal in terms ofits individual pixel values and in terms of Fourier modes., Thus they are a compromise between characterising a signal in terms of its individual pixel values and in terms of Fourier modes. + In the first case. the characterisation has no information on correlations between dillerent pixels. (no. frequency information) but perfect. positional information.," In the first case, the characterisation has no information on correlations between different pixels (no frequency information) but perfect positional information." + A Fourier decomposition. on the other hand. has perfect. frecucney information but no positional information.," A Fourier decomposition, on the other hand, has perfect frequency information but no positional information." + The analysis of a spectrum in terms of wavelets has the advantage that one can study the clustering of lines (positional information)). as à function of their widths frequency )).," The analysis of a spectrum in terms of wavelets has the advantage that one can study the clustering of lines positional ), as a function of their widths frequency )." + The usage of wavelets to analyse QSO spectra was xoneered by Pando lang (19096. 1998). who used a wavelet analysis of absorption lines to describe the clustering of those ines.," The usage of wavelets to analyse QSO spectra was pioneered by Pando Fang (1996, 1998), who used a wavelet analysis of absorption lines to describe the clustering of those lines." + The wavelet analysis detected. large. scale. structure in the forest. which had. proved dillicult using more racditional methods.," The wavelet analysis detected large scale structure in the forest, which had proved difficult using more traditional methods." + In contrast to Pando Fane. we will use wavelets to analyse the absorption spectrum. directly. hereby. eliminating the somewhat subjective step of first decomposing the continuous spectrum in absorption lines.," In contrast to Pando Fang, we will use wavelets to analyse the absorption spectrum directly, thereby eliminating the somewhat subjective step of first decomposing the continuous spectrum in absorption lines." + The advantage of this new method. is that it allows. us o objectively characterise the typical width of absorption eatures as a function of position along the., The advantage of this new method is that it allows us to objectively characterise the typical width of absorption features as a function of position along the. +. We will show using hyerodyvnamic simulations that the obability distribution of wavelet amplitudes can be used to characterise the equation of state of the absorbing medium. in terms of the temperature at the mean density. Z5. and the slope. >. ofthe temperature-density relation.," We will show using hydrodynamic simulations that the probability distribution of wavelet amplitudes can be used to characterise the equation of state of the absorbing medium, in terms of the temperature at the mean density, $T_0$, and the slope, $\gamma$, of the temperature-density relation." + In addition we use the fact that wavelets are localised in position along the spectrum. thereby. allowing us to detect. spatial variations in Z5 and/or 7. which might be present as a result of late helium LE reionization or local effects.," In addition we use the fact that wavelets are localised in position along the spectrum, thereby allowing us to detect spatial variations in $T_0$ and/or $\gamma$, which might be present as a result of late helium II reionization or local effects." + This paper is organised as follows., This paper is organised as follows. + In Section 2. we firs eive a brief description of the generation of mock spectra from our simulations and illustrate the decomposition of the spectra in discrete wavelets., In Section \ref{sect:setup} we first give a brief description of the generation of mock spectra from our simulations and illustrate the decomposition of the spectra in discrete wavelets. + The statistics of the wavele amplitudes for dillerent simulations is cliscussecl in Section 3 and the results are summarised. in Section 4.., The statistics of the wavelet amplitudes for different simulations is discussed in Section \ref{sect:analysis} and the results are summarised in Section \ref{sect:conclusions}. + Recently. Aleiksin (2000) discussed. indepentlv the application of wavelets to QSO spectra.," Recently, Meiksin (2000) discussed indepently the application of wavelets to QSO spectra." + We use the LI simulation described before in Theuns (2000), We use the L1 simulation described before in Theuns (2000). +" 3riclly. this is a simulation of a flat. vacuum energv dominated: cold. dark matter model with matter density Q,,;=0.3. barvon fraction Ομ=049 and Llubble constant ο=65 km ! "," Briefly, this is a simulation of a flat, vacuum energy dominated cold dark matter model with matter density $\Omega_m=0.3$, baryon fraction $\Omega_b h^2=0.019$ and Hubble constant $H_0=65$ km $^{-1}$ $^{-1}$." +Density Ductuations in this model are normalised to the abundance of galaxy clusters (Eke 1996). and we have used. (Seljak Zaldarriaga 1996) to compute the appropriate linear. transfer function., Density fluctuations in this model are normalised to the abundance of galaxy clusters (Eke 1996) and we have used (Seljak Zaldarriaga 1996) to compute the appropriate linear transfer function. + The ICM in this model is photo-jionised. ancl photo-heated by the UV-background. from QSOs. as computed by LIaardt Macau (1996).," The IGM in this model is photo-ionised and photo-heated by the UV-background from QSOs, as computed by Haardt Madau (1996)." + We simulated this cosmological model with a mocified version of the simulation code (Couchman 1995). which combines hierarchical. P3M. gravity (Couchman 1991) with smoothed particle hvdrodynamics (SPL. Lucy 1977. Gingold Monaghan 1977).," We simulated this cosmological model with a modified version of the simulation code (Couchman 1995), which combines hierarchical P3M gravity (Couchman 1991) with smoothed particle hydrodynamics (SPH, Lucy 1977, Gingold Monaghan 1977)." + We simulate a periodic. cubie box of size 7.7 co-moving Alpe using 128? articles of cach species. which gives us sullicient resolution o compute line-widths reliably (Pheuns 1905].," We simulate a periodic, cubic box of size 7.7 co-moving Mpc using $^3$ particles of each species, which gives us sufficient resolution to compute line-widths reliably (Theuns 1998)." + To investigate other cllects. we also make use of simulations of a model with the same numerical resolution. Cosmology ancl hermal history. but with a smaller box size (3.8 Alpe). and a set of simulations with a smaller normalisation ex=0.775 and ox= 0.4.," To investigate other effects, we also make use of simulations of a model with the same numerical resolution, cosmology and thermal history, but with a smaller box size (3.8 Mpc), and a set of simulations with a smaller normalisation $\sigma_8=0.775$ and $\sigma_8=0.4$ ." + In the analysis stage. we impose a particular equation of state on the gas at. low overdensities (ορ)« 20) of the form T=Tolp/ip}) varving the values of Zi and 5.," In the analysis stage, we impose a particular equation of state on the gas at low overdensities $\rho/\langle\rho\rangle < 20$ ) of the form $T=T_0 (\rho/\langle\rho\rangle)^{\gamma-1}$, varying the values of $T_0$ and $\gamma$." + We then compute mock spectra that mimick the actual observed LILIES spectrum of the sa.=3.0 QSO L107|485. discussed. by Rauch et al. (," We then compute mock spectra that mimick the actual observed HIRES spectrum of the $z_{\rm em}=3.0$ QSO 1107+485, discussed by Rauch et al. (" +1997). using the following procedure.,"1997), using the following procedure." + We divide the observed. spectrum in three redshifts bins. z=252625. ;=2.6252.875 and z=2.8753 and scale the mean absorption of the simulations at 2=2.5. 2.=2.75 and 2=3 to the corresponding observed value.," We divide the observed spectrum in three redshifts bins, $z=2.5-2.625$, $z=2.625-2.875$ and $z=2.875-3$ and scale the mean absorption of the simulations at $z=2.5$, $z=2.75$ and $z=3$ to the corresponding observed value." + “Phe simulated. spectra are resampled to the observed. resolution. and. convolved. with a Gaussian to mimick instrumental broadening.," The simulated spectra are resampled to the observed resolution, and convolved with a Gaussian to mimick instrumental broadening." + We have analysed the noise statistics of the QSO 1107 spectrum as a function of Lux. ancl add: noise with these properties tothe simulated: spectra.," We have analysed the noise statistics of the QSO 1107 spectrum as a function of flux, and add noise with these properties tothe simulated spectra." + By randomly. combining individual sight lines through the simulation volume. we generate a," By randomly combining individual sight lines through the simulation volume, we generate a" +using à mocificd version of the program:VCUBINE.,using a modified version of the program. +. ΔΙΣ used to fit the strong component. ενος)., M2 used to fit the strong component $_{strong}$ ). + The orbital parameters determined from for both methocls are shown in Table 2.., The orbital parameters determined from for both methods are shown in Table \ref{table2}. + Phe velocity values from vanWaelkens.&Waters(1995) are also presented in ‘Table 2.., The velocity values from \citet{vanwinckel1995} are also presented in Table \ref{table2}. + The quantities below the horizontal rule (ay. Mr. and m») in Table 2 are derived using the elective inclination angle of 35° and à primary mass of 0.5 M.," The quantities below the horizontal rule $_{1}$ , $_{T}$, and $_{2}$ ) in Table \ref{table2} are derived using the effective inclination angle of $\degr$ and a primary mass of 0.8 $_{\odot}$." + Aside from the orbital period. difference. the orbital parameters found. via A2 are in very close agreement with those of vanWinckel.Waelkens.&Waters (1995).," Aside from the orbital period difference, the orbital parameters found via M2 are in very close agreement with those of \citet{vanwinckel1995}." +. This close agreement suggests hat the RY measurements using M2 produce parameters for he svstem are more consistent with Waelkensetal.(1996)., This close agreement suggests that the RV measurements using M2 produce parameters for the system are more consistent with \citet{waelkens1996}. +. 1t should be noted that the uncertainty in the mass function and secondary mass are most sensitive to the uncertainty in he semi-amplitucle., It should be noted that the uncertainty in the mass function and secondary mass are most sensitive to the uncertainty in the semi-amplitude. + Dased ont10 orbital data determined using the it to the strong component (Νεο). a to-scale schematic of the binary orbits has been constructed. to aid. in. the interpretation of the spectra. see Fig. 2..," Based on the orbital data determined using the fit to the strong component $_{strong}$ ), a to-scale schematic of the binary orbits has been constructed to aid in the interpretation of the spectra, see Fig. \ref{fig2}." + The locations of the 33 APO observations to date are also indicated in the figure along with kev orbital locations., The locations of the 33 APO observations to date are also indicated in the figure along with key orbital locations. + Orbital phase zero. Ó—0. was chosen to correspond with the location of phase zero defined by Waelkensetal.(1996). to aid in comparison οἱ data.," Orbital phase zero, $\phi = 0$, was chosen to correspond with the location of phase zero defined by \citet{waelkens1996} to aid in comparison of data." + Waelkensetal.(1996) chose 6=0 to correspond to.D) 2448300. near minimum RY.," \citet{waelkens1996} chose $\phi = 0$ to correspond to JD 2448300, near minimum RV." + Since the components of the binary are in an eccentric orbi. the Roche lobe radius is not constant as a function of orbial phase.," Since the components of the binary are in an eccentric orbit, the Roche lobe radius is not constant as a function of orbital phase." + Phe consequences of the variable size of the Roc1¢ lobe are discussed in 5.3.., The consequences of the variable size of the Roche lobe are discussed in \ref{sec-phase_dependent-asymmetries}. + In. panel a of Fig., In panel a of Fig. + 10 the toche lobe radius of the primary star was calculated as a function of orbital phase using equation (46) of Sepinsky.Wilems.&Ixalogera (2007)., \ref{fig10} the Roche lobe radius of the primary star was calculated as a function of orbital phase using equation (46) of \citet{sepinsky2007}. +". Phe exact value of the Roc16 racius may οου, as will be discussed later:however. the relative variation and orbital phase dependence should remain the same."," The exact value of the Roche radius may differ, as will be discussed later;however, the relative variation and orbital phase dependence should remain the same." +"In order to determine the absolute magnitudes of the galaxies observed in these fields, K corrections must be applied for an accurate estimate.","In order to determine the absolute magnitudes of the galaxies observed in these fields, $K$ corrections must be applied for an accurate estimate." +" We have determined i-band K corrections for these galaxies assuming both a starburst type galaxy and an Sa type galaxy, with the templates originating from the SDSS galaxy spectra template libraries available at/index."," We have determined $i$ -band $K$ corrections for these galaxies assuming both a starburst type galaxy and an Sa type galaxy, with the templates originating from the SDSS galaxy spectra template libraries available at." +"html.. Although in principle the galaxies hosting the absorbers in question could be of early type, the fact that there is significant continuum emission at short wavelengths apparent from the g band images and flat spectral energy distributions for these galaxies indicates that star formation is currently ongoing in these systems."," Although in principle the galaxies hosting the absorbers in question could be of early type, the fact that there is significant continuum emission at short wavelengths apparent from the $g$ band images and flat spectral energy distributions for these galaxies indicates that star formation is currently ongoing in these systems." +" Hence, we have chosen to include only the templates for these star forming galaxies for the K corrections."," Hence, we have chosen to include only the templates for these star forming galaxies for the $K$ corrections." + Assuming an early type galaxy spectrum would increase the absolute magnitudes of these galaxies., Assuming an early type galaxy spectrum would increase the absolute magnitudes of these galaxies. +" The K corrected absolute magnitudes are given in Table 4.1,, assuming that each galaxy is at the redshift of the absorber."," The $K$ corrected absolute magnitudes are given in Table \ref{Tab:Results}, assuming that each galaxy is at the redshift of the absorber." + The value of M;=—21.59 from Blanton(2003) was used for determining luminosities., The value of $_{i}^{\star}=-21.59$ from \citet{Blanton03} was used for determining luminosities. +" In total, the fields of Q1009—0026, Q1228+1018, Q1323—0021, and Q1436—0051 all contain L>L* galaxies within 75 kpc projected impact parameters."," In total, the fields of $-$ 0026, Q1228+1018, $-$ 0021, and $-$ 0051 all contain $L>L^{\star}$ galaxies within 75 kpc projected impact parameters." + Could these galaxies be unrelated field galaxies that just happen to lie near the sightline to the QSOs?, Could these galaxies be unrelated field galaxies that just happen to lie near the sightline to the QSOs? +" This is a persistent question when imaging QSO absorbers, which can ultimately only be settled with followup spectroscopy to confirm the redshifts of the galaxies."," This is a persistent question when imaging QSO absorbers, which can ultimately only be settled with followup spectroscopy to confirm the redshifts of the galaxies." +" However, the probability of observing"," However, the probability of observing" +"where m are magnitudes, σ=0.921x10-9""g, σ are photometric uncertainties, and the sum includes N pairs verifying v;—r;~Ar.","where $m$ are magnitudes, $\overline{\sigma} = 0.921 \times 10^{- 0.4 m} \sigma$, $\sigma$ are photometric uncertainties, and the sum includes $N$ pairs verifying $\tau_j - \tau_i \sim +\Delta \tau$." +" Here, L=109"" are monochromatic luminosities in convenient units and r are restframe times (e.g.,CidFernandesetal.2000;Goicoechea 2008)."," Here, $L = 10^{- 0.4 m}$ are monochromatic luminosities in convenient units and $\tau$ are restframe times \citep[e.g.,][]{Cid00,Goi08}." +. This SF(L) describes typical luminosity variabilities at different restframe lags., This $SF(L)$ describes typical luminosity variabilities at different restframe lags. +" We normalize the original structure function to the luminosity variance and then take the square-root for convenience, i.e., we analyse the normalized structure function f=ISFOI?JoQD)."," We normalize the original structure function to the luminosity variance and then take the square-root for convenience, i.e., we analyse the normalized structure function $f = [SF(L)]^{1/2} / \sigma(L)$." + Restframe lags substantially below the restframe duration of the records are considered in the analysis., Restframe lags substantially below the restframe duration of the records are considered in the analysis. +" In the initial selection, we take ArxP/4(1+ z), where P is the duration of each whole combined record (see Fig. 1))"," In the initial selection, we take $\Delta \tau \leq P/4(1 + z)$ , where $P$ is the duration of each whole combined record (see Fig. \ref{combi}) )." +" Later, only lags before reaching the asymptotic behaviours (f< 1) are taken into account."," Later, only lags before reaching the asymptotic behaviours $f \leq$ 1) are taken into account." +" In Fig. 2,,"," In Fig. \ref{sfgr}," +" using LQLM I observations, the normalized structure function at A~ 2100 ((g band; filled circles) is compared to f at A~ 2600 ((r band; open circles)."," using LQLM I observations, the normalized structure function at $\lambda \sim$ 2100 $g$ band; filled circles) is compared to $f$ at $\lambda \sim$ 2600 $r$ band; open circles)." + Both growths seem to be consistent with each other., Both growths seem to be consistent with each other. +" With respect to the initial logarithmic slopes, in Fig."," With respect to the initial logarithmic slopes, in Fig." + 2 we also show two fits f=A(Arf with f~7x/Naof values close to 1 (Nao is the number of degrees of freedom)., \ref{sfgr} we also show two fits $f = A (\Delta \tau)^{\beta}$ with $\hat{\chi}^2 = \chi^2/N_{dof}$ values close to 1 $N_{dof}$ is the number of degrees of freedom). +" The fits over time intervals 3--14 d (g band) and 6-50 d (r band) lead to 8 = 0.71 + 0.08 and B = 0.63 + 0.02 (1o intervals), respectively."," The fits over time intervals $-$ 14 d $g$ band) and $-$ 50 d $r$ band) lead to $\beta$ = 0.71 $\pm$ 0.08 and $\beta$ = 0.63 $\pm$ 0.02 $\sigma$ intervals), respectively." +" These initial slopes disagree with the prediction of the cellular-automaton disc-instability model, but roughly agree with the time-symmetric flares that appear in the one-dimensional hydrodynamical simulations (see Introduction)."," These initial slopes disagree with the prediction of the cellular-automaton disc-instability model, but roughly agree with the time-symmetric flares that appear in the one-dimensional hydrodynamical simulations (see Introduction)." +" In principle, the starburst model could also account for the measured slope at ~ 2100 ((g band)."," In principle, the starburst model could also account for the measured slope at $\sim$ 2100 $g$ band)." +" The microlensing slope 6~ 0.2-0.3 (Hawkins2002) is strongly inconsistent with the LQLM I data ofQ0957+561,, which is not at all surprising."," The microlensing slope $\beta +\sim$ $-$ 0.3 \citep{Haw02} is strongly inconsistent with the LQLM I data of, which is not at all surprising." +" We are studying intrinsic fluctuations, so microlensing does not play any role."," We are studying intrinsic fluctuations, so microlensing does not play any role." +" In Fig. 3,,"," In Fig. \ref{sfgg}," + the APO (open triangles) and LQLM I (filled circles) structure functions of at ~ 2100 hhave significantly different initial growths., the APO (open triangles) and LQLM I (filled circles) structure functions of at $\sim$ 2100 have significantly different initial growths. +" This may be a consequence of evolution in the variability scenario, since both experiments (APO and LQLM I) are separated by ~ 1500 d in the quasar restframe."," This may be a consequence of evolution in the variability scenario, since both experiments (APO and LQLM I) are separated by $\sim$ 1500 d in the quasar restframe." +" In fact, the oldest (APO) brightness records seem to incorporate relatively short fluctuations (which would not be present in the LQLM I light curves) that would generate the differences between initial growths."," In fact, the oldest (APO) brightness records seem to incorporate relatively short fluctuations (which would not be present in the LQLM I light curves) that would generate the differences between initial growths." + The solid lines in Fig., The solid lines in Fig. + 3 fit the APO behaviours in three time intervals ($?~ 1)., \ref{sfgg} fit the APO behaviours in three time intervals $\hat{\chi}^2 \sim$ 1). +" Their slopes are8 = 0.81 + 0.12 at Arx 5 d,B = 0.66 + 0.06 at Ar€ 11 d, andB = 0.59 + 0.04 at At< 16d (1c intervals)."," Their slopes are $\beta$ = 0.81 $\pm$ 0.12 at $\Delta \tau \leq$ 5 d, $\beta$ = 0.66 $\pm$ 0.06 at $\Delta \tau \leq$ 11 d, and $\beta$ = 0.59 $\pm$ 0.04 at $\Delta \tau \leq$ 16 d $\sigma$ intervals)." +" In order to discuss the mechanism(s) for (M)UV variability, and to quantify thespectral and time evolution of that mechanism(s), we compare the"," In order to discuss the mechanism(s) for (M)UV variability, and to quantify thespectral and time evolution of that mechanism(s), we compare the" +"select about 300 directions in the 2?x60? simulated sky, each of which goes through at least one galaxy-scale halo with a mass above 10?!Mo located at redshift |z—0.6|«0.02 in the replicated box.","select about 300 directions in the $2^{\circ} \times 60^{\circ}$ simulated sky, each of which goes through at least one galaxy-scale halo with a mass above $10^{12}h^{-1}M_{\odot}$ located at redshift $|z-0.6| \leqslant 0.02$ in the replicated box." + This ensures that the primary haloes we select are responsible for producing multiple images of the Zs=2 background sources., This ensures that the primary haloes we select are responsible for producing multiple images of the $z_s=2$ background sources. +" We have confirmed that these ~ 300 randomly selected primary lenses are representative of mass and circular velocity distributions of ~23,000 haloes that meet the same selection criteria in the simulated sky."," We have confirmed that these $\sim$ 300 randomly selected primary lenses are representative of mass and circular velocity distributions of $\sim$ 23,000 haloes that meet the same selection criteria in the simulated sky." +" As stated in §44.2, velocity dispersions of the simulated main lensing haloes range from 200kms to 300kms~’, comparable to our sample of observed lenses! (see the end of §33.1)."," As stated in 4.2, velocity dispersions of the simulated main lensing haloes range from $\kms$ to $\kms$, comparable to our sample of observed lenses (see the end of 3.1)." +" As the selection function of the observed lenses is hard to define, we assume that they are a random sample of haloes in the ranges of velocity dispersion and redshift given above."," As the selection function of the observed lenses is hard to define, we assume that they are a random sample of haloes in the ranges of velocity dispersion and redshift given above." +" Furthermore, by imposing a lower mass limit of 10?5!Me on the main lens, we have excluded cases where two or more less massive haloes aligned along the same line of sight produce a comparable strong-lensing signal."," Furthermore, by imposing a lower mass limit of $10^{12}h^{-1}M_{\odot}$ on the main lens, we have excluded cases where two or more less massive haloes aligned along the same line of sight produce a comparable strong-lensing signal." +" However, the chance is small for two foreground haloes, both more massive than 10'°h~!Mo, to be well aligned to jointly lens a background galaxy."," However, the chance is small for two foreground haloes, both more massive than $10^{10}h^{-1}M_{\odot}$, to be well aligned to jointly lens a background galaxy." +" Emprically, such “three-dimensional” lenses are rare: in the CLASS survey, only one probable such case has been reported out of a total of 22 candidates (??))."," Emprically, such ``three-dimensional'' lenses are rare: in the CLASS survey, only one probable such case has been reported out of a total of 22 candidates \citealt{CMA01B2114,Augusto01B2114}) )." + Fig., Fig. + 8 shows the geometry and halo distribution of an example lensing cone., \ref{fig:LCOA50} shows the geometry and halo distribution of an example lensing cone. + All haloes within a given cone are used for the lensing calculation., All haloes within a given cone are used for the lensing calculation. +" On average each lensing cone (of 50""x50"") contains about 10,000 (12,000) haloes (subhaloes)."," On average each lensing cone (of $50\arcsec\times +50\arcsec$ ) contains about $10,000$ $12,000$ ) haloes (subhaloes)." +" Within a projected central region of R«5"" for strong lensing, there are on average ~300 haloes with πι>108h-1Mo directly contributing to the convergence field."," Within a projected central region of $R\leqslant +5\arcsec$ for strong lensing, there are on average $\sim300$ haloes with $m>10^8h^{-1}M_{\odot}$ directly contributing to the convergence field." + The rest are distributed further out (in projection) and contribute to the shear field of this region in the same way as point masses., The rest are distributed further out (in projection) and contribute to the shear field of this region in the same way as point masses. +" To carry out calculations for multi-plane light deflection, we assume 60 lens planes distributed with equal spacing in redshift between the observer and the source at z,=2."," To carry out calculations for multi-plane light deflection, we assume 60 lens planes distributed with equal spacing in redshift between the observer and the source at $z_s=2$." +" In each of these lens planes, a region of 5”x around the line centre is covered by a 1000x rigid grid in order to calculate the Jacobian matrix Ας (Eq. 2))"," In each of these lens planes, a region of $5\arcsec\times5\arcsec$ around the line centre is covered by a $1000\times1000$ rigid grid in order to calculate the Jacobian matrix $A_{\rm s}$ (Eq. \ref{eq:Jacobian}) )" + between the source plane and the final image plane., between the source plane and the final image plane. + Haloes within a lensing cone are projected into these lens planes according to their redshifts., Haloes within a lensing cone are projected into these lens planes according to their redshifts. +" The main lens halo is modelled as an isothermal ellipsoid, for which a universal axis ratio (qs= 0.8) and core radius (so=0.05"") are assumed."," The main lens halo is modelled as an isothermal ellipsoid, for which a universal axis ratio $q_3=0.8$ ) and core radius $s_0=0.05\arcsec$ ) are assumed." +" The orientation of the ellipsoid is randomly chosen in the interval of [0, 27]."," The orientation of the ellipsoid is randomly chosen in the interval of [0, $\pi$ ]." +" The lensing strength bsrg (related with br through br=bsrme/sin~+e, where e=(1—q3)'/?, see ?)) is derived through an empirical relationship between halo’s virial velocity Vooo and the velocity dispersion ostg of the equivalent isothermal ellipsoid (?)): and Όσιο=4m(cosig/c)?Da./D,., where c is the speed of light and Das and D, are the angular diameter distances between the main lens and the source, and the source and the observer, respectively."," The lensing strength $b_{\rm SIE}$ (related with $b_I$ through $b_I=b_{\rm + SIE}e/sin^{-1}e$, where $e=(1-q_3^2)^{1/2}$ , see \citealt{KKprofile1998}) ) is derived through an empirical relationship between halo's virial velocity $V_{200}$ and the velocity dispersion $\sigma_{\rm SIE}$ of the equivalent isothermal ellipsoid \citealt{ChaeMao2006}) ): and $b_{\rm SIE}=4\pi(\sigma_{\rm SIE}/c)^2D_{\rm ds}/D_{\rm s}$, where $c$ is the speed of light and $D_{\rm ds}$ and $D_{\rm s}$ are the angular diameter distances between the main lens and the source, and the source and the observer, respectively." + The virial velocity Vooo is obtained from halo mass Mooo and its virial radius rooo through Vào=GMo00/r200., The virial velocity $V_{200}$ is obtained from halo mass $M_{200}$ and its virial radius $r_{200}$ through $V_{200}^2=GM_{200}/r_{200}$. +" Our requirement that the main lens be more massive than 10138!Mc results in osre ranging from 200kms~' to 300kms~', with a weighted mean (by cross-section, ος o?) of kms""! corresponding to a lensing strength of bsm=0.84"" for our adopted lens and source redshifts (z4=0.6,zs 2.0)."," Our requirement that the main lens be more massive than $10^{12} h^{-1}M_{\odot}$ results in $\sigma_{\rm SIE}$ ranging from $\kms$ to $\kms$, with a weighted mean (by cross-section, $\propto \sigma^4$ ) of $\kms$ corresponding to a lensing strength of $b_{\rm SIE}=0.84\arcsec$ for our adopted lens and source redshifts $z_d=0.6, z_s=2.0$ )." +" Within each lensing cone, haloes with projected profiles that are completely outside the central 5”x5"" region are treated as point masses."," Within each lensing cone, haloes with projected profiles that are completely outside the central $5\arcsec\times5\arcsec$ region are treated as point masses." +" Those within this region are assigned a density profile: as described above, we investigate three distinct choices of this profile (singular isothermal sphere, NFW with the M08 concentration-mass relation, and NFW with the B01-MO05 concentration-mass relation)."," Those within this region are assigned a density profile: as described above, we investigate three distinct choices of this profile (singular isothermal sphere, NFW with the M08 concentration-mass relation, and NFW with the B01-M05 concentration-mass relation)." + All halo profiles are normalized to their masses Maoo and truncated at the virial radii r290; subhaloes are truncated at two times their half mass radii., All halo profiles are normalized to their masses $M_{200}$ and truncated at the virial radii $r_{200}$; subhaloes are truncated at two times their half mass radii. +" For each line of sight, deflection angles are individually calculated for the equivalent isothermal ellipsoid of the main lens and for all line-of-sight (sub)haloes, and are tabulated to the meshesat different lens planes."," For each line of sight, deflection angles are individually calculated for the equivalent isothermal ellipsoid of the main lens and for all line-of-sight (sub)haloes, and are tabulated to the meshesat different lens planes." +" Through ray tracing, source positions Bw that correspond to the final image plane δι are identified, and the final Jacobian matrix"," Through ray tracing, source positions $\vec{\beta}_N$ that correspond to the final image plane $\vec{\theta}_1$ are identified, and the final Jacobian matrix" +debris perturbed by the LAIC alone.”,debris perturbed by the LMC alone.” + If applied to the 100 kpe NW MOI stream. the tenth Fourier terii would IC a 0-367 yvariation. which is roughly comparable to a binning where there still is real structure relative to a constant. but only on vet larger scales.," If applied to the 100 kpc NW M31 stream, the tenth Fourier term would be a $\phi$ variation, which is roughly comparable to a binning where there still is real structure relative to a constant, but only on yet larger scales." +" Majewskietal.(2001). considered kincmatic data for he Sagittarius stream in our galaxy. aud conchided that he halo could not be very hup. however we note hat our sinulatious show that sub-hialos tend to ""chop"" streams leaving locally cold reiunauts with pieces at the sale radius offset in velocity. which may be compatible with the Majewskietal.(2001). Figure 2."," \citet{Majewski:04} considered kinematic data for the Sagittarius stream in our galaxy and concluded that the halo could not be very lumpy, however we note that our simulations show that sub-halos tend to “chop"" streams leaving locally cold remnants with pieces at the same radius offset in velocity, which may be compatible with the \citet{Majewski:04} Figure 2." + Kiucinatic tests will become very powerful iu testing for substructure ouce large velocity samples are available., Kinematic tests will become very powerful in testing for substructure once large velocity samples are available. + The 100 kpe NW stream in MOI is a coherent geometric structure whose orbit is well clear of the body of M31 naking it a near ideal testing eround for the presence or absence of thousands of dark sub-halos predicted imu LCDAM ii-body simulations., The 100 kpc NW stream in M31 is a coherent geometric structure whose orbit is well clear of the body of M31 making it a near ideal testing ground for the presence or absence of thousands of dark sub-halos predicted in LCDM n-body simulations. + Although the stream appears o be nearly a half ellipse over a conuuon range of uctallicities the upper brauch of the stream is less well defined and has a mumber of clearly visible gaps., Although the stream appears to be nearly a half ellipse over a common range of metallicities the upper branch of the stream is less well defined and has a number of clearly visible gaps. + The lower branch is που] complete and provides a uuch more conservative test for the presence of sub-iilos., The lower branch is nearly complete and provides a much more conservative test for the presence of sub-halos. + The main result of this paper is that the stream was highly significant density variations on virtually all scales from 2 kpe. up to about 20 kpc.," The main result of this paper is that the stream has highly significant density variations on virtually all scales from 2 kpc, up to about 20 kpc." + The variation of density around the mean has a very low probability of beine a chance statistical fluctuation. less than 10.," The variation of density around the mean has a very low probability of being a chance statistical fluctuation, less than $10^{-5}$." + We lave been careful to measure the huupiness relative to an averaged local background and masked out gaps and regious where the photometric uncertaimties add significantly to the variations., We have been careful to measure the lumpiness relative to an averaged local background and masked out gaps and regions where the photometric uncertainties add significantly to the variations. + As a control suuple we take a higher metallicity set of stars at exactly the same location as the stream. which finds no significant deusity variations.," As a control sample we take a higher metallicity set of stars at exactly the same location as the stream, which finds no significant density variations." + Relative to other known cool star streams. the ADT stream stauds apart for its length. distance from the disturbing effects of the MOI disk. a variety of scales of substructure and its hieh statistical significance.," Relative to other known cool star streams, the M31 stream stands apart for its length, distance from the disturbing effects of the M31 disk, a variety of scales of substructure and its high statistical significance." + It is interesting to uote that M31 appears to have oulv the one well defined long stellar stream at larec radius., It is interesting to note that M31 appears to have only the one well defined long stellar stream at large radius. + The NW= stream has quite a low total Iuninosity so it should not particularly stand out relative to other stellar streams that were created at large radius over the buildup of the M21 dark halo., The NW stream has quite a low total luminosity so it should not particularly stand out relative to other stellar streams that were created at large radius over the buildup of the M31 dark halo. + Several other stream fraeimieuts are visible im the full field map Richardsonetal.(2011). but none as long or coherent as tle NW stream., Several other stream fragments are visible in the full field map \citet{Richardson:11} but none as long or coherent as the NW stream. + [Tf oue asstuned that other loug streams likely formed over the lifetime of the halo then the absence of others at the present time indirectly sugeests that they have been broken up by sub-halos (Carlbere2009)., If one assumed that other long streams likely formed over the lifetime of the halo then the absence of others at the present time indirectly suggests that they have been broken up by sub-halos \citep{StarStreams}. +. The measured NW ΕΠ ΠΠ essentially rules out the possibilitv that it is a low mass star stream orbiting in a smooth galactic potential of a disk plus bulee plus dark halo., The measured NW stream lumpiness essentially rules out the possibility that it is a low mass star stream orbiting in a smooth galactic potential of a disk plus bulge plus dark halo. + Couverselv. the NW stream deusity variations are compatible with the level of deusitv changes that à laree population of dark matter sub-lhialos induce.," Conversely, the NW stream density variations are compatible with the level of density changes that a large population of dark matter sub-halos induce." + The details of the density variation are scusitive to the statistical distribution of the sub-structure but are unlikely to be specifically modeled for a single stream., The details of the density variation are sensitive to the statistical distribution of the sub-structure but are unlikely to be specifically modeled for a single stream. + As more streams are observed statistical modcling of the degree of substructure preseut should become possible., As more streams are observed statistical modeling of the degree of substructure present should become possible. + The present study is statistically consistent with a hiehlv sub-structured dark halo. although we caution that the density data alone for a single stream is nof a conclusive proof that halos are as sub-structured as LCDAL simulations predict.," The present study is statistically consistent with a highly sub-structured dark halo, although we caution that the density data alone for a single stream is not a conclusive proof that halos are as sub-structured as LCDM simulations predict." + Overall. the density variations of the NW= stream are strong circunistautial evidence that the predicted thousands. of dark matter sub-halos are present iu M231's dark halo.," Overall, the density variations of the NW stream are strong circumstantial evidence that the predicted thousands of dark matter sub-halos are present in M31's dark halo." + This research is supported by NSERC. CIfAR. aud NRC in Canada and the French ANR programe POAMAINE.," This research is supported by NSERC, CIfAR and NRC in Canada and the French ANR programme POMMME." +crossings.,crossings. + The cusp event parameters shows a peak width lower-limit corresponding to the caustic crossing correlation., The cusp event parameters shows a peak width lower-limit corresponding to the caustic crossing correlation. + Vherefore. the systematic bias introduced. into the source size determination by the assumption that the 1988 event was due to a caustic crossing can only result in an over estimate of source size.," Therefore, the systematic bias introduced into the source size determination by the assumption that the 1988 event was due to a caustic crossing can only result in an over estimate of source size." + In. addition. Fig.," In addition, Fig." + 9. demonstrates that a cusp event contains very little information on source size/transverse velocity due to the lack of any correlation with peak height., \ref{height_v_width} demonstrates that a cusp event contains very little information on source size/transverse velocity due to the lack of any correlation with peak height. + The inference that the 1988. peak was a caustic crossing and the 19909 peak was a cusp event is consistent with Fig. 9.., The inference that the 1988 peak was a caustic crossing and the 1999 peak was a cusp event is consistent with Fig. \ref{height_v_width}. . + Fig., Fig. + 2 shows a AA.~0.8 magnitude rise between the 1997 image C minimum and the 1998 level. suggesting an event in between those observing seasons.," \ref{data} shows a $\Delta M_{obs}\sim 0.8$ magnitude rise between the 1997 image C minimum and the 1998 level, suggesting an event in between those observing seasons." + We assume that the intrinsic source luminosity was approximately. constant over this period. which is supported both bv. the [acts that image A changed by <0.2 magnitudes and that the other images show opposite trends.," We assume that the intrinsic source luminosity was approximately constant over this period, which is supported both by the facts that image A changed by $\la0.2$ magnitudes and that the other images show opposite trends." + Comparing the image € change to the probabilities in Fig., Comparing the image C change to the probabilities in Fig. + 6 we find that ANAL. is Consistent with a [|ee caustic crossing having occurred between the 1997 and 1998 observing seasons. but rules out re caustic crossings ( 99%) and cusp events (~955 ," \ref{assym} we find that $\Delta M_{obs}$ is consistent with a $+ve$ caustic crossing having occurred between the 1997 and 1998 observing seasons, but rules out $-ve$ caustic crossings $\sim 99\%$ ) and cusp events $\sim 95\%$ )." +We therefore infer that a [|ee caustic crossing was missed between the LOOT and 1998 observing seasons., We therefore infer that a $+ve$ caustic crossing was missed between the 1997 and 1998 observing seasons. + In this section we calculate. probability functions for the likelihood of observing future LEMES in images A and € given current light-curves., In this section we calculate probability functions for the likelihood of observing future HMEs in images A and C given current light-curves. + In this section we assume that there was à [|re caustic crossing between the LOOT ancl 1998 observing seasons. ancl investigate when we should next see a caustic crossing in image €. These calculations follow Witt. Ixavser Itefscal (1993) who caleulated the separations in cdimensionless units of the cdillerent combinations of |re and οὐ events.," In this section we assume that there was a $+ve$ caustic crossing between the 1997 and 1998 observing seasons, and investigate when we should next see a caustic crossing in image C. These calculations follow Witt, Kayser Refsdal (1993) who calculated the separations in dimensionless units of the different combinations of $+ve$ and $-ve$ events." + llowever we have included both our estimates of n? and Dopp. and the cusp as a third class of LEAL.," However we have included both our estimates of $\langle m\rangle$ and $v_{eff}$, and the cusp as a third class of HME." + Due to the typical cüamond formation of fold caustics. the case of à. fre followed. by ai re caustic crossing is common.," Due to the typical diamond formation of fold caustics, the case of a $+ve$ followed by a $-ve$ caustic crossing is common." + Similarly. inspection of mocel light-curves shows that cusp events follow —ee caustic crossings as the source moves past the cusp associated. with that caustic his feature is seen in the double horned profile that is characteristic of the Chang-Iefsdal lens).," Similarly, inspection of model light-curves shows that cusp events follow $-ve$ caustic crossings as the source moves past the cusp associated with that caustic (this feature is seen in the double horned profile that is characteristic of the Chang-Refsdal lens)." + However. we have inferred that the OGLE image € light-curve shows a [re caustic crossing followed by a cusp event., However we have inferred that the OGLE image C light-curve shows a $+ve$ caustic crossing followed by a cusp event. + Such a combination is much less common anc is due to the source moving past a cusp formed from two fold causties other than the one responsible for the |ee caustic crossing HELME., Such a combination is much less common and is due to the source moving past a cusp formed from two fold caustics other than the one responsible for the $+ve$ caustic crossing HME. + 1 can also be seen in mocdel light-curves., It can also be seen in model light-curves. + Examples of the two scenarios are shown in Fig. 10.., Examples of the two scenarios are shown in Fig. \ref{scenarios_real}. + The upper panel shows an example of a double horned. event., The upper panel shows an example of a double horned event. + Phe lower two panels show examples of a double horned event surrounding a cusp event., The lower two panels show examples of a double horned event surrounding a cusp event. + The source is shown passing very close to the cusp (centre panel). partly coming in contact. ancl passing further away. producing a lower amplitude event (lower panel).," The source is shown passing very close to the cusp (centre panel), partly coming in contact, and passing further away, producing a lower amplitude event (lower panel)." + The ieht-curves were produced. using our most likely mocel for he microlensing parameters (epp=300msee1. (n)=LIAL. and 8—5.10! em).," The light-curves were produced using our most likely model for the microlensing parameters $v_{eff}=300km\,sec^{-1}$, $\langle m\rangle=0.1\,M_{\odot}$ and $S=5\times10^{14}\,cm$ )." + The intervals (og., The intervals (eg. + 1-2) quoted xlow refer to the intervals between the events labelled on hese plots., 1-2) quoted below refer to the intervals between the events labelled on these plots. + From these two scenarios. we generate the probability unctions for 4 different event separation statistics (plotted in Fig. 11)).," From these two scenarios, we generate the probability functions for 4 different event separation statistics (plotted in Fig. \ref{caust_prob}) )." + The left handand. right hand. panels show, The left handand right hand panels show +for the zero-metal protostars.,for the zero-metal protostars. + An interesting result from the Yoshii Saio calculations is that the IME derived would have à maximum around the intermediate stellar mass range (3-8 M. )., An interesting result from the Yoshii Saio calculations is that the IMF derived would have a maximum around the intermediate stellar mass range (3-8 $_\odot$ ). + Verv recently. Nakamura&Umemnura(2000). performed. multi-dimensional hvdrodynamie simulations of the collapse aud fragmentation of filamentary primordial clouds.," Very recently, \citet{nak00} performed multi-dimensional hydrodynamic simulations of the collapse and fragmentation of filamentary primordial clouds." + These simulations show that. depending upon the initial clensitw of the cloud. the IME for Population HI stars is likely to be bimodal.," These simulations show that, depending upon the initial density of the cloud, the IMF for Population III stars is likely to be bimodal." + Gas filaments with initial densities lower than ~10? * tend to fragment into structures with masses larger than several tens of AL.(e107 ML. ). while initially denser filaments (7>10° 7) experiment more effective II cooling and fragment into structures of 1—2 M...," Gas filaments with initial densities lower than $\sim 10^5$ $^{-3}$ tend to fragment into structures with masses larger than several tens of $M_\odot$$\sim 10^2$ $_\odot$ ), while initially denser filaments $n\ga 10^5$ $^{-3}$ ) experiment more effective $_2$ cooling and fragment into structures of $\sim 1-2$ $_\odot$." + The relative peaks of this bimodal IME would be a funcüon of the collapse epoch in such a wav (hat during the first epoch the dominant peak would be around 2 M... while as the star formation proceeds and the collapsing clump heats up. the Jeans mass would be displaced to larger masses. moving the peak to ~LO? M...," The relative peaks of this bimodal IMF would be a function of the collapse epoch in such a way that during the first epoch the dominant peak would be around $2$ $_\odot$, while as the star formation proceeds and the collapsing clump heats up, the Jeans mass would be displaced to larger masses, moving the peak to $\sim 10^2$ $_\odot$." + The above mentioned studies basically compile the proposed IMESs for Population I] stars in the literature., The above mentioned studies basically compile the proposed IMFs for Population III stars in the literature. + Ii summary: 1) there are plivsical arguments for assuming that the slope of the present IAIF has changed in Gime. i) the first star-lormine clouds ivpically fragmented into massive clumps of LO?—10* M... iii) there exists the possibility of further fragmentation into smaller masses. producing an IMF peaking at about the intermediate stellar mass range.," In summary: i) there are physical arguments for assuming that the slope of the present IMF has changed in time, ii) the first star-forming clouds typically fragmented into massive clumps of $10^2-10^3$ $_\odot$, iii) there exists the possibility of further fragmentation into smaller masses, producing an IMF peaking at about the intermediate stellar mass range." + Support to these [acts hasbeen recently eiven by Hernandez Ferrara (2001)., Support to these facts hasbeen recently given by Hernandez Ferrara (2001). + In the framework of the standard. hierarchical clustering scenario of galaxy formation. these authors show that the IMIF of the first stars was increasingly high mass biassed towards high redshifts: at 2~9 the charactheristic stellar mass being 10-15 AL..," In the framework of the standard hierarchical clustering scenario of galaxy formation, these authors show that the IMF of the first stars was increasingly high mass biassed towards high redshifts: at $z\sim 9$ the charactheristic stellar mass being 10-15 $_\odot$." + Therefore. we will restrict our analysis to the IMFESs that cover the above properties.," Therefore, we will restrict our analysis to the IMFs that cover the above properties." + In fact. we use the IMESs proposed by Yoshii Saio. relerring to them as YSa if a relation mass-huninosity in the form L(m)xGnM.) with 4=1.5 is adopted. or YSb if 3= 3.," In fact, we use the IMFs proposed by Yoshii Saio, referring to them as YSa if a relation mass-luminosity in the form $L(m)\propto (m/M_\odot)^\beta$ with $\beta=1.5$ is adopted, or YSb if $\beta=3$ ." + We also use the bimodal IME proposed by Nakamura Umenmura ancl. in a similar way. we," We also use the bimodal IMF proposed by Nakamura Umemura and, in a similar way, we" +"δη Ry= eesimtilerM co3021) Note that while Va, is a coutribution to the variance about the mean Aj. it is nof necessarily noise!"," G(R) = r (kr)W(r) Note that while $\Var_{\rm 3pt}$ is a contribution to the variance about the mean $\Delta_j$, it is not necessarily noise!" + Mich of it the density around the particular object. which of course has scatter frou the mean.," Much of it the density around the particular object, which of course has scatter from the mean." + The shot noise or Poisson contribution to the variance is based onu the expected counts. includiug clustering. which ave ΕΠ.ιδ].," The shot noise or Poisson contribution to the variance is based on the expected counts, including clustering, which are $\nbar[1+w_{is}(R)]$." + Tho variance is D. the homogenous shot noise becomes 2677ΕΕ097.," The variance is _j) = ]), the homogenous shot noise becomes $2a^2\nbar/(V\phi_0)^2$." + Figure 1. shows the contribution per radial yn to the variance in the shot noise., Figure \ref{fig:g} shows the contribution per radial bin to the variance in the shot noise. + Esseutially all of the shot noise arises at Rx15a: in other words. the fact that oue is subtracting the backeround with a region at moderate radius rather than the eutire saze adds little extra noise to the density estimator.," Essentially all of the shot noise arises at $R<1.5a$; in other words, the fact that one is subtracting the background with a region at moderate radius rather than the entire sample adds little extra noise to the density estimator." +" The four above Vois. Vai. homogeneous shot noise, and clustered shot noiseall have different scalings with the «epth of the survey. the size of the window. aud the clustering strenethl."," The four above $\Var_{\rm 2pt}$, $\Var_{\rm 3pt}$, homogeneous shot noise, and clustered shot noise—all have different scalings with the depth of the survey, the size of the window, and the clustering strength." + For a typical survey thickuess Lfog and a typical window radiusαν the contributions to the variance are roughly (L/200)A. AP (ES)354)(1/a?09). aud (A/32)(17/009). respectively.," For a typical survey thickness $L\equiv \nbar/\phi_0$ and a typical window radius$a$ , the contributions to the variance are roughly $(L/20a)\Delta$, $\Delta^2$, $(L/125a)(1/a^3\phi_0)$, and $(\Delta/32)(1/a^3\phi_0)$, respectively." + The wmmerical coefficieuts are for illustrationoulvl., The numerical coefficients are for illustration. +. The 2-poiut custering fermi dominates on large scales: the clustered. shot-noise on πα scales., The 2-point clustering term dominates on large scales; the clustered shot-noise on small scales. + The process of averagingeoO many A;] iuto a mean A iutroduces additional error terus from the correlations of the spectroscopic galaxy posiions. including coutributious from the four-point correlation fiuction.," The process of averaging many $\Delta_j$ into a mean $\Delta$ introduces additional error terms from the correlations of the spectroscopic galaxy positions, including contributions from the four-point correlation function." +" This is not surprising becase Aisa uiuteeral of e,CR). whose covariance normally imvolves the four-»omt function."," This is not surprising because $\Delta$ is an integral of $\wis(R)$, whose covariance normally involves the four-point function." + Neelecting these new terms 1 favor of the equations above corresponds to the assurtion that the shot noise of the spectroscopic sample exceeds its clusteriug (seetheexpansionsofDerusteiu1991:Tanulton100}.," Neglecting these new terms in favor of the equations above corresponds to the assumption that the shot noise of the spectroscopic sample exceeds its clustering \citep[see the +expansions of][]{Ber94,Ham97}." + This is often a gouc assulupion on small scales. particularly if oue is considering oily a sia] subset of the spectroscopic sanupe.," This is often a good assumption on small scales, particularly if one is considering only a small subset of the spectroscopic sample." + A crucial assuuption of the aualvsis is that the objects that are uncorrelate but iucose projection with the spectroscoxc objec are statistically ideutical to those in other parts of the skv., A crucial assumption of the analysis is that the objects that are uncorrelated but in close projection with the spectroscopic object are statistically identical to those in other parts of the sky. + Magification from weak leusimecau violate this assiunuption iu principle (es...Valottoetal.1997).," Magnification from weak lensingcan violate this assumption in principle \citep[e.g.,][]{Val97}." +.. Aoreover. selection biases in cluster catalogs owing to superposition of unrelated structures (c.g..Valottoeta.2001). are not reduced by this inethod.," Moreover, selection biases in cluster catalogs owing to superposition of unrelated structures \citep[e.g.,][]{Val01} are not reduced by this method." + If we are given GR). wecau use the Abel integral to find the corresponding Wr):," If we are given $G(R)$ , wecan use the Abel integral to find the corresponding $W(r)$ :" +"used by Rochauetal.(2010) and the observations of Brandletal.(1999) sample the core region of the cluster, where the density of cluster stars is highest, we conclude that the in the external regions from the models of Robinetal.(2003) is reasonable and might actually be an overestimate of the true value.","used by \cite{roc10} and the observations of \cite{bra99} sample the core region of the cluster, where the density of cluster stars is highest, we conclude that the in the external regions from the models of \cite{rob03} is reasonable and might actually be an overestimate of the true value." + We therefore assume it as an upper limit to the contamination level in this field., We therefore assume it as an upper limit to the contamination level in this field. +" We will return in refsecjetttotheworko f Rochauetal. (2010)and Brandletal.(1999)todi ise85 dola icationsforthestudyo fthecluster' sstetarpopulation, bi "," We will return in \\ref{sec_lett} to the work of \cite{roc10} + and \cite{bra99} to discuss their implications for the study of the cluster's stellar population, but we first need to address the issue of differential extinction in this field." +An efficient method to quantify the amount of differential extinction in a star cluster is to use the position of a star in the observed CMD to calculate its distance from a fiducial line (e.g. the ZAMS itself) along the direction of the reddening vector., An efficient method to quantify the amount of differential extinction in a star cluster is to use the position of a star in the observed CMD to calculate its distance from a fiducial line (e.g. the ZAMS itself) along the direction of the reddening vector. +" This distance would be the resultant of two components, namely E(V-I) on the abscissa and Ay on the ordinate, and would give us an estimate of the extinction toward the star itself (seeanexampleinPiottoetal.1999)."," This distance would be the resultant of two components, namely E(V-I) on the abscissa and $_V$ on the ordinate, and would give us an estimate of the extinction toward the star itself \citep[see an example in][]{pio99}." +". This method works under the assumption that the observed stars are at the same distance, i.e. belong to the same system."," This method works under the assumption that the observed stars are at the same distance, i.e. belong to the same system." +" As discussed above, the field star contamination in the range of magnitudes covered by our data is high."," As discussed above, the field star contamination in the range of magnitudes covered by our data is high." +" The described method can be applied in a reliable way only to bright objects (V« 17), since these are very likely cluster members and field star contamination is minimized at these magnitudes."," The described method can be applied in a reliable way only to bright objects $<17$ ), since these are very likely cluster members and field star contamination is minimized at these magnitudes." +" Although our data cannot be used for this purpose, since all stars brighter than V~17.5 are saturated, using shorter exposures SB04 were able to perform a study of differential reddening in 33603 taking advantage of multi-band HST photometry of the bright massive stars (i.e. the same objects that are saturated in our images)."," Although our data cannot be used for this purpose, since all stars brighter than $V\sim17.5$ are saturated, using shorter exposures SB04 were able to perform a study of differential reddening in 3603 taking advantage of multi-band HST photometry of the bright massive stars (i.e. the same objects that are saturated in our images)." + These authors were able to map the variation of E(B-V) as a function of the distance from the cluster centre (see their Figure 5b)., These authors were able to map the variation of E(B-V) as a function of the distance from the cluster centre (see their Figure 5b). +" They found that the value Ay~4.5 is representative of the very centre of the OB stars association, while they noticed an increase toward the external regions."," They found that the value $_V \simeq 4.5$ is representative of the very centre of the OB stars association, while they noticed an increase toward the external regions." +" Following the work of SB04 and adopting RyΞ3.55 as they suggest, we estimate that the mean value of Ay in the area our observations (from ~10"" to ~ 70"") is Ay-5.5."," Following the work of SB04 and adopting $R_V =3.55$ as they suggest, we estimate that the mean value of $_V$ in the area sampled by our observations (from $\sim10\arcsec$ to $\sim70\arcsec$ ) is $_V=5.5$." +" sampledTherefore,by in the rest of this paper we will adopt this value to correct our magnitudes for extinction."," Therefore, in the rest of this paper we will adopt this value to correct our magnitudes for extinction." + The final CMD corrected in this way is shown in Figure 3.., The final CMD corrected in this way is shown in Figure \ref{fig_cmd_dered}. +" Once extinction is taken into account, the CMD shown in Figure 3 can be used to determine stellar ages through PMS isochrone fitting."," Once extinction is taken into account, the CMD shown in Figure \ref{fig_cmd_dered} can be used to determine stellar ages through PMS isochrone fitting." +" PMS isochrones with ages of 1, 2, 3, 10, 20 and 30 Myr from Siessetal.(2000) are shown in the figure from to left."," PMS isochrones with ages of 1, 2, 3, 10, 20 and 30 Myr from \cite{sie00} are shown in the figure from right to left." + We used the same value of the distance modulus rightadopted for the ZAMS fit., We used the same value of the distance modulus adopted for the ZAMS fit. + Note that the Siessetal.(2000) models are not available for the WFC3 photometric system., Note that the \citet{sie00} models are not available for the WFC3 photometric system. +" As above, we used the ATLASO library of Kurucz(1993) to calculate the magnitude differences between the JC to WFC3 photometric systems."," As above, we used the ATLAS9 library of \cite{ku93} to calculate the magnitude differences between the JC to WFC3 photometric systems." +" After this correction, the models indicate that the lowest mass that we reach for PMS stars is 0.3 Mo."," After this correction, the models indicate that the lowest mass that we reach for PMS stars is 0.3 $_{\odot}$." +" In the magnitude range V«20, where the photometric uncertainty on the V—7 color (~ 0.05) is smaller than the typical isochrone separation, the CMD suggests for our PMS stars an age in the range from 1 to 10 Myr, with an average age of 3 Myr."," In the magnitude range $V<20$, where the photometric uncertainty on the $V-I$ color $\sim 0.05$ ) is smaller than the typical isochrone separation, the CMD suggests for our PMS stars an age in the range from 1 to 10 Myr, with an average age of 3 Myr." +" Assuming the PMS isochrones are correct, this can already be considered as tentative evidence of an age spread in the stellar population of 33603."," Assuming the PMS isochrones are correct, this can already be considered as tentative evidence of an age spread in the stellar population of 3603." +" HAO8 studied the IMF of NGC33603 using near-infraredRecently, (IR) imaging from ground-based adaptive optics photometry of the cluster center obtained with the NAOS-CONICA (NACO) camera and the wider field Infrared Spectrometer and Array Camera (ISAAC) at the Very Large Telescope (VLT)."," Recently, HA08 studied the IMF of 3603 using near-infrared (IR) imaging from ground-based adaptive optics photometry of the cluster center obtained with the NAOS--CONICA (NACO) camera and the wider field Infrared Spectrometer and Array Camera (ISAAC) at the Very Large Telescope (VLT)." +" Their JHKL’ bands photometry reaches the magnitude limit J20.5 (ie. ~0.4M for cluster stars) in an area of —110"" radius from the cluster center.", Their $JHKL'$ bands photometry reaches the magnitude limit $J\sim20.5$ (i.e. $\sim 0.4$ $_{\odot}$ for cluster stars) in an area of $\sim 110\arcsec$ radius from the cluster center. +" By comparing their CMD to Baraffeetal.(1998) PMS evolutionary models, they identify a population of PMS stars with ages of 0.5—1.0 MMyr."," By comparing their CMD to \cite{ba98} PMS evolutionary models, they identify a population of PMS stars with ages of $0.5-1.0$ Myr." +" Moreover, by adopting a set of MS isochrones from Lejeune&Schaerer(2001), they provide tentative hints of the presence of an evolved MS population of 2.02.5 MMyr."," Moreover, by adopting a set of MS isochrones from \cite{le01}, they provide tentative hints of the presence of an evolved MS population of $2.0-2.5$ Myr." + Note that this age estimate is based on the comparison- of the MS isochrones with the position of three massive evolved O stars in their IR CMDs., Note that this age estimate is based on the comparison of the MS isochrones with the position of three massive evolved O stars in their IR CMDs. +" Coupling this piece of evidence with the presence of the evolved post-red supergiant star 225 (Moffat1983) in the cluster field, HAO8 hypothesize a possible age spread in the cluster population suggesting the presence of two distinct bursts in the star formation history, separated by ~10 MMyr."," Coupling this piece of evidence with the presence of the evolved post-red supergiant star 25 \citep{mof83} in the cluster field, HA08 hypothesize a possible age spread in the cluster population suggesting the presence of two distinct bursts in the star formation history, separated by $\sim 10$ Myr." +" It is interesting to note that recently Melenaetal.(2008) have Sher225 at the same distance as NGC 3603, suggesting a placedcommon origin."," It is interesting to note that recently \cite{mel08} have placed 25 at the same distance as NGC 3603, suggesting a common origin." +" The same ISAAC observations analyzed by HAO08, were previously used by Stolteetal.(2004) and Brandletal.(1999) to study the low-mass stars population in 33603."," The same ISAAC observations analyzed by HA08, were previously used by \cite{sto04} and \cite{bra99} to study the low-mass stars population in 3603." +" While by inspecting the CMDs shown in their Figure 4 and 3c, respectively, signatures of the presence of a low MS population in the CMDs can not be ruled out, both these papers agree on dating a PMS population in the range of 0.5-1 MMyr."," While by inspecting the CMDs shown in their Figure 4 and 3c, respectively, signatures of the presence of a low MS population in the CMDs can not be ruled out, both these papers agree on dating a PMS population in the range of $0.5-1$ Myr." + SB04 published what was by then the deepest optical CMDs based on a combination of UBVI and Ha photometry from the Siding Spring Observatory (SSO) and archival HST/WFPC2 observations., SB04 published what was by then the deepest optical CMDs based on a combination of $UBVI$ and $H{\alpha}$ photometry from the Siding Spring Observatory (SSO) and archival HST/WFPC2 observations. + The SSO ground-based CMDs sample the brightest massive population., The SSO ground-based CMDs sample the brightest massive population. + The WFPC2 CMD, The WFPC2 CMD +5.5 ~ I.. 2.,$\times$ $\times$ $\sim$ \ref{fig:rawimg}. +" 2). O0PLE10.5. 30°22""00° OOPLEM20.5° 30722?457. 2.. ~2.1/", \ref{fig:smoimg} \ref{fig:smoimg}) $00\hd14\md19.5\sd$ $-30^\circ24\md09\sd$ $00\hd14\md20.5\sd$ $-30^\circ22\md45\sd$ \ref{fig:smoimg}. +" ;;-model ~640 ή 1.0. 00L4""18/7. 30723""!167. :7-model. ro=(474+33)3=0.79+0.06 (lo ?-model ((~900 (O2)... 1.7+0.2 Clo) "," $\sim$ $\beta$ $\sim$ $\beta$ $1.0$ $00\hd14\md18.7\sd$ $ -30^\circ23\md16\sd$ $\beta$ $r_c = (474 \pm 33)$$\beta = +0.79 \pm 0.06$ $\sigma$ $\beta$ $\sim$ $1.7\pm0.2$ $\sigma$ " +SSP-equivaleut age of the blue pixel sequence at Vo20 mag aresec7 rauges from 5 Mr (at Voe16 mag 2) to 300 Myr (at V—20 Dias 7: with solar metallicity and se=1 asstuuptious).,SSP-equivalent age of the blue pixel sequence at $V<20$ mag $^{-2}$ ranges from 5 Myr (at $V\sim16$ mag $^{-2}$ ) to 300 Myr (at $V\sim20$ mag $^{-2}$; with solar metallicity and $\tau_V=1$ assumptions). + This pixel age estimation agrees with the result of Twang&Lee(2010) that the star cluster formation rate in M51 increased sienificautly duriug the period of 100250 Myr ago., This pixel age estimation agrees with the result of \citet{hwa10} that the star cluster formation rate in M51 increased significantly during the period of $100-250$ Myr ago. + Since the pixels along spiral arius tend to be relatively bright (V-2021 nag D2j we infer that the brightest tip of the blue pixel sequence inav be brighter for a ealaxy with brighter aud more obvious spiral arius.," Since the pixels along spiral arms tend to be relatively bright $V<20-21$ mag $^{-2}$ ), we infer that the brightest tip of the blue pixel sequence may be brighter for a galaxy with brighter and more obvious spiral arms." + However. to use the brightest tip of the blue pixel sequence as a star formation rate mdicator. it should be considered that the brightest tip stronglv depends on spatial resolution. as shown in Fig.," However, to use the brightest tip of the blue pixel sequence as a star formation rate indicator, it should be considered that the brightest tip strongly depends on spatial resolution, as shown in Fig." + 8 aud Tables 3-. 1., \ref{pcmdbin} and Tables \ref{pcmdvar1}- \ref{pcmdvar2}. + In addition to the brightest tip. the color dispersion at given surface brightness may be also a possible parameter reflecting the star formation historv of a galaxy.," In addition to the brightest tip, the color dispersion at given surface brightness may be also a possible parameter reflecting the star formation history of a galaxy." + Lauvon-Fosteretal.(2007) suggested that total blue/red light ratio of a galaxy can be used as a 1norphloloev iudicator for that galaxy. showing a good correlation between the total blueτος ratio and galaxy morphological tvpe.," \citet{lan07} suggested that total blue/red light ratio of a galaxy can be used as a morphology indicator for that galaxy, showing a good correlation between the total blue/red ratio and galaxy morphological type." + As an extension of such an idea. we investigated the blue/red ratio as a function of pixel surface brightuess in refcdist..," As an extension of such an idea, we investigated the blue/red ratio as a function of pixel surface brightness in \\ref{cdist}." + As a result. we found that the pCMD of NGC 5191 shows a remarkae feature of the due/red ratio peak at 19.5xV20.0. mae’ 7. which is consistently found even iu he low-resolution versions of he pCMD up to 100« ànung (200 pe | resolution).," As a result, we found that the pCMD of NGC 5194 shows a remarkable feature of the blue/red ratio peak at $19.5\le V<20.0$ mag $^{-2}$, which is consistently found even in the low-resolution versions of the pCMD up to $100\times100$ binning (200 pc $^{-1}$ resolution)." + Tuterestinely. the peak blue/red ratio itself iu he BV pCMD remains consistent within c uncertainty up to the 50« binning (100 ος l yesolutiou) pCMD (0.532+ 0.061).," Interestingly, the peak blue/red ratio itself in the $B-V$ pCMD remains consistent within $\sigma$ uncertainty up to the $50\times50$ binning (100 pc $^{-1}$ resolution) pCMD $0.532\pm0.061$ )." +" ILlowever. the peak blue/red ratio in the V1 X""MD ends to decrease as the binning factor Dicreases."," However, the peak blue/red ratio in the $V-I$ pCMD tends to decrease as the binning factor increases." + The pixels around the surface brightuess for the ue]fred ratioB peak peak~E19.7— mag D 7) reveal the spiral-zxin patterus well. indicating that he blue/red ratio iu this range is a measure of the spiral avin fraction iu this galaxy.," The pixels around the surface brightness for the blue/red ratio peak $\mu_{peak}\sim19.7$ mag $^{-2}$ ) reveal the spiral-arm patterns well, indicating that the blue/red ratio in this range is a measure of the spiral arm fraction in this galaxy." + This parameter nay be a morphology indicator better than the total blue/red lieht ratio. because it is derived in a coutrolled surface brightucss range.," This parameter may be a morphology indicator better than the total blue/red light ratio, because it is derived in a controlled surface brightness range." + The fipeak does not depeud on spatial resolution. whe PAoiial resolution is finer than ~LOO pe |.," The $\mu_{peak}$ does not depend on spatial resolution, when spatial resolution is finer than $\sim100$ pc $^{-1}$." + Tlowever. it needs to be checked whether pipes 38 universal or depending on individual properties of spiral galaxies.," However, it needs to be checked whether $\mu_{peak}$ is universal or depending on individual properties of spiral galaxies." +" Iu principle. the total color itself of a galaxy cau rot be a perfect morphology indicator. because were are solae unusual types of galaxics such as blue early-type galaxies and τος late-type ealaxies (c.e,Leeetal.2008.2010a.b.c)."," In principle, the total color itself of a galaxy can not be a perfect morphology indicator, because there are some unusual types of galaxies such as blue early-type galaxies and red late-type galaxies \citep[e.g.][]{lee08,lee10a,lee10b,lee10c}." +. However. 1ο conibinatioun of Ενω aud the blue/red ratio ανα mav be useful to discriminate hic carly-type galaxies from usual late-type galaxies. jecause blue early-type galaxies are known to be duest at their brightest center (thatis.thosurfaceLeeetal.2006. 2008).," However, the combination of $\mu_{peak}$ and the blue/red ratio at $\mu_{peak}$ may be useful to discriminate blue early-type galaxies from usual late-type galaxies, because blue early-type galaxies are known to be bluest at their brightest center \citep[that is, the surface brightness distributions of their blue pixels may be different from those of usual late-type galaxies;][]{lee06,lee08}." +". This needs to be tested in ""ture studies.", This needs to be tested in future studies. + The spatial distributions of pixels iu different color domains show how stellar populations vary according to their location., The spatial distributions of pixels in different color domains show how stellar populations vary according to their location. + In NGC 5101. there seeni to be two important factors affecting the internal distribution of stellar populations.," In NGC 5194, there seem to be two important factors affecting the internal distribution of stellar populations." + The first is the effect of spiral arias., The first is the effect of spiral arms. + In Fig. 16..," In Fig. \ref{iring}," + from the *before-aruy area to the cafter-aruy area. it is clear that the fraction of the vouug population (Dl) increases lavecly (by factor of ~1.9 or 2.6) while the vouug aud dusty population (P2) decreases (by factor of ~0.6 or 0:8).," from the `before-arm' area to the `after-arm' area, it is clear that the fraction of the young population (P1) increases largely (by factor of $\sim1.9$ or 2.6) while the young and dusty population (P2) decreases (by factor of $\sim0.6$ or 0.8)." + These results are explained by the density wave model (Lindblad1963:Tooumwe1977:Bertinetal. 1989).. in which railing spiral deusitv waves proceed slower than rotating material. compressing eas and dust into 1ος stars.," These results are explained by the density wave model \citep{lin63,too77,ber89}, in which trailing spiral density waves proceed slower than rotating material, compressing gas and dust into new stars." + The spatial distributions of pixels across he spiral axis of NGC 5191 are consistent with such a compressing process: deuse dust (before-arni) > newly-formed stars (aftter-arm)., The spatial distributions of pixels across the spiral arms of NGC 5194 are consistent with such a compressing process: dense dust (before-arm) $\rightarrow$ newly-formed stars (after-arm). + The jxel« highh-extinct by dust result in the plume catures in the pCMD., The pixels highly-extinct by dust result in the plume features in the pCMD. + Another factor is the effect of the interaction with NGC 5195., Another factor is the effect of the interaction with NGC 5195. + While the pixels iu the imuer disk of NGC 5191 have relatively svinmetric properties (that is. the trends aloug Al-A2-A3 are similar to those along QaALA5S-AG). the pixels iu he outer," While the pixels in the inner disk of NGC 5194 have relatively symmetric properties (that is, the trends along A1-A2-A3 are similar to those along A4-A5-A6), the pixels in the outer" +NGOC6397. region hot enough for some additional Li burning to occur.,"NGC6397, region hot enough for some additional Li burning to occur." + The trend. of lithium abundance with effective temperature along the sub-eiant branch. found. in M4 by Mucciarellietal.(2011) is similar to that found by Gonzalezllernándezetal.(2000). in NGC6397," The trend of lithium abundance with effective temperature along the sub-giant branch, found in M4 by \citet{muc} is similar to that found by \citet{jonay} in NGC6397." + This brief summary of previous investigations shows how problematic. from a theoretical point of view. it is to determine the primordial Li abundance of stars.," This brief summary of previous investigations shows how problematic, from a theoretical point of view, it is to determine the primordial Li abundance of stars." + In this paper we investigate a complementary avenue to estimate the Lalo primordial A(Li)., In this paper we investigate a complementary avenue to estimate the Halo primordial A(Li). + We employ spectroscopy of Halo stars evolving along the lower RGB. defined Following Grattonetal.(2000) as the portion of the RGB brighter than the luminosity corresponding o the completion of the first. dredge-up and fainter than he ItGDB-bump.," We employ spectroscopy of Halo stars evolving along the lower RGB, defined – following \citet{gratton00} – as the portion of the RGB brighter than the luminosity corresponding to the completion of the first dredge-up and fainter than the RGB-bump." +. We will show that the ellect of atomic diffusion on the surface Li abundances of these objects is much smaller than lor stars., We will show that the effect of atomic diffusion on the surface Li abundances of these objects is much smaller than for stars. + As a consequence. he observed. A(Li) can be emploved. to set. independent. very strong constraints on any additional physical. process (ic. turbulent mixing. preprocessing of Llalo material. mocifications to the BBN) eventually needed to reconcile hese values with BBN predictions.," As a consequence, the observed A(Li) can be employed to set independent, very strong constraints on any additional physical process (i.e., turbulent mixing, preprocessing of Halo material, modifications to the BBN) eventually needed to reconcile these values with BBN predictions." + The paper is structured. as follows., The paper is structured as follows. + Section 2 analyzes he theoretical. advantages of cmploving (Li) measured in the atmosphere of lower RGB Lalo stars to estimate heir primordial Li abundance. while Sect.," Section 2 analyzes the theoretical advantages of employing A(Li) measured in the atmosphere of lower RGB Halo stars to estimate their primordial Li abundance, while Sect." + 3 discusses our selected: sample of lower RGB halo objects and the derivation of A(Li) and. Fe/LH]., 3 discusses our selected sample of lower RGB halo objects and the derivation of A(Li) and [Fe/H]. + Section 4 presents theoretical oedietions for the surface Li abundance in Population 11 ower RGB stars from mocels including atomic cliffusion and convection as element transport mechanisms — and the derivation of the initial Li for our selected star sample. that oovides an estimate of the Halo primordial Li abundance independent of thePlateau.," Section 4 presents theoretical predictions for the surface Li abundance in Population II lower RGB stars – from models including atomic diffusion and convection as element transport mechanisms – and the derivation of the initial Li for our selected star sample, that provides an estimate of the Halo primordial Li abundance independent of the." + Section 5 discusses the Li abundance in the lower RGB of the Galactic globulars GC 6397. NGC 6752 and MA.," Section 5 discusses the Li abundance in the lower RGB of the Galactic globulars NGC 6397, NGC 6752 and M4." + A summary and conclusions close the paper., A summary and conclusions close the paper. + The surface Li abundance after the first. dredge-up is essentially a consequence of the dilution due to the increased size of the convective envelope after the MS turn oll (plus a minor contribution from Li burning in the deep lavers of the fully mixed envelope in the most metal rich Halo stars), The surface Li abundance after the first dredge-up is essentially a consequence of the dilution due to the increased size of the convective envelope after the MS turn off (plus a minor contribution from Li burning in the deep layers of the fully mixed envelope in the most metal rich Halo stars). +" At the end of the MS phase. when the deepening convective boundary. reaches lavers where the Li-burning (15,5 72.5 10"" IX) was ellicient during. the MS. the surface A(Li) begins to decrease."," At the end of the MS phase, when the deepening convective boundary reaches layers where the Li-burning $T_{burn}\sim$ 2.5 $10^6$ K) was efficient during the MS, the surface A(Li) begins to decrease." + Εις depletion essentially ends when the convective envelope attains its maximum depth and the first dredge-up is complete., This depletion essentially ends when the convective envelope attains its maximum depth and the first dredge-up is complete. + An important point to notice is that atomic diffusion during the ALS produces a local maximum of the Li abundance in the radiative lavers right below the convective envelope (see.c.g..Richarcetal.2005).. for only a relatively small fraction of the envelope Li is transported deep enough to reach lavers where it is eventually burned.," An important point to notice is that atomic diffusion during the MS produces a local maximum of the Li abundance in the radiative layers right below the convective envelope \citep[see, e.g.,][]{richard05}, for only a relatively small fraction of the envelope Li is transported deep enough to reach layers where it is eventually burned." + As a consequence. models without anc with ciffusion fully cllicient or even moderated. by levitation or some turbulence that just mixes back material cilfused from the envelope — cülute similar amount of Li within the deepening convective region.," As a consequence, models without and with diffusion -- fully efficient or even moderated by levitation or some turbulence that just mixes back material diffused from the envelope – dilute similar amount of Li within the deepening convective region." + In addition. the maximum size of the convective envelope is also weakly alleeted by. clilfusion this can be also inferred. by the fact that the predicted RGB bump luminosities with anc without diffusion are very similar (Cassisietal.L997:Michaud:2010) and the resulting (Li) abundances on the lower HOB are only slightly changed.," In addition, the maximum size of the convective envelope is also weakly affected by diffusion -- this can be also inferred by the fact that the predicted RGB bump luminosities with and without diffusion are very similar \citep{cas97, mic10} – and the resulting A(Li) abundances on the lower RGB are only slightly changed." +" ""This is in contrast with models. [or upper AIS stars. where the elect of diffusion on the surface abundances can reach several tenths of dex (seeforexam-pleFig.3inMucciarellietal."," This is in contrast with models for upper MS stars, where the effect of diffusion on the surface abundances can reach several tenths of dex \citep[see for example Fig.~3 in][]{muc}." +2011).. Theoretical models also predict that along the RGB. after the completion of the dredge-up. atomic dilfusion and levitation are not able to modify. appreciably the surface abundance of Li and all other elements (Michaudοἱal.2007.2010).," Theoretical models also predict that along the RGB, after the completion of the dredge-up, atomic diffusion and levitation are not able to modify appreciably the surface abundance of Li and all other elements \citep{mic07, mic10}." +. The chemical abundance measurements in giant stars by Grattonetal.(2000)... Spiteetal.(2005). and. Lind.οἱal.(2009).. sugσισον that any additional element. transport along the RGB is very likely inellicient in this phase.," The chemical abundance measurements in giant stars by \citet{gratton00}, \citet{spite05} and \citet{lind}, suggest that any additional element transport along the RGB is very likely inefficient in this phase." + On the other hand. stars evolved: bevond the RGB bump clisplay the clleet of an additional mixing event. for which the so-called. thermohaline mixing! is nowadays the most. popular candidate (see.e.g.Charbonnel&Lagarde2010.andrefer-encestherein )..," On the other hand, stars evolved beyond the RGB bump display the effect of an additional mixing event, for which the so-called 'thermohaline mixing' is nowadays the most popular candidate \citep[see, e.g.][and references therein]{charb}." + These considerations lead to the conclusion that. in general there is a simpler relationship between initial ancl current surface Li abundances in lower RGB stars compared toPlateau objects., These considerations lead to the conclusion that in general there is a simpler relationship between initial and current surface Li abundances in lower RGB stars compared to objects. + The basic reason is hat. for a fixed initial Li abundance. lower RGB surface abundances are sensitive to the total amount of Li left in he star.," The basic reason is that, for a fixed initial Li abundance, lower RGB surface abundances are sensitive to the total amount of Li left in the star." + In fact. after the effect of dilution is accounted or. they are. allected only by the total amount of Li »irned. during the ALS phase due to atomic cdillusion. plus »xossible additional clement transport mechanisms — if they are invoked to solve the discrepancy with BBN results.," In fact, after the effect of dilution is accounted for, they are affected only by the total amount of Li burned during the MS phase due to atomic diffusion plus possible additional element transport mechanisms – if they are invoked to solve the discrepancy with BBN results." + On he other hand. the observed. abundances along the are determined by the evolution of the rate of Li depletion from the convective envelope due just to dilfusion or moderatedenhanced. by additional turbulent processes.," On the other hand, the observed abundances along the are determined by the evolution of the rate of Li depletion from the convective envelope due just to diffusion or moderated/enhanced by additional turbulent processes." + Combining the information obtained from Li in lower BD stars with data. sets much stronger constraints on how these proposed additional mixing processes act between the base of the convective envelope and the deeper Li burning regions.," Combining the information obtained from Li in lower RGB stars with data, sets much stronger constraints on how these proposed additional mixing processes act between the base of the convective envelope and the deeper Li burning regions." + In this study. we address the issue of the Lalo primordial Li abundance by coupling predictions from standard. ROB stellar models to the Li abundances measured on a sample of lower ROB Population LL field stars.," In this study, we address the issue of the Halo primordial Li abundance by coupling predictions from standard RGB stellar models to the Li abundances measured on a sample of lower RGB Population II field stars." + As ciscussed before. this will provide an independent. quantitative estimate of the dillerence (i£. anv) with the Bie Bane value. ancl set a very robust constraint on the elliciency of additional physical processes invoked. το resolve. this. discrepancy. complementary to the interpretation of the abundances along thePlateau.," As discussed before, this will provide an independent quantitative estimate of the difference (if any) with the Big Bang value, and set a very robust constraint on the efficiency of additional physical processes invoked to resolve this discrepancy, complementary to the interpretation of the abundances along the." +" Wo the solution of the discrepancy involves pre-processing curing the Galaxy. formation or mocifications to the BBN (sec.e.g.Piauetal.2006:Cvburtetal. 2010).. our analysis will provide a solid quantitative estimate of. respectively. the net amount ofLi burned during the carly Galaxy evolution. or exactly how much Li must be synthesized in the revised"" BBN."," If the solution of the discrepancy involves pre-processing during the Galaxy formation or modifications to the BBN \citep[see, e.g.,][]{piau06, cyburt10}, our analysis will provide a solid quantitative estimate of, respectively, the net amount of Li burned during the early Galaxy evolution, or exactly how much Li must be synthesized in the 'revised' BBN." + To the best of our. knowledge. analyses of the Li abundances measured. alone the RGB have. been so. [ar mainly aimed at testing the agreement between observed and. predicted. Li depletion. after the first. dredge-up.," To the best of our knowledge, analyses of the Li abundances measured along the RGB have been so far mainly aimed at testing the agreement between observed and predicted Li depletion after the first dredge-up," +This is also the time lag between the peaks (uot the beeiuuings) of svuchrotrou flares at ro and vy. provided that the helt-crossine time is uniniportant relative to the electron injecting time. aud that >mar (C8. Zhaug et al.,"This is also the time lag between the peaks (not the beginnings) of synchrotron flares at $\nu_2$ and $\nu_1$, provided that the light-crossing time is unimportant relative to the electron injecting time, and that $\gamma \ll \gamma_{max}$ (e.g., Zhang et al." + 2002: Chiappoetti et al., 2002; Chiappetti et al. + 1999: Chiaberge (κοπή 1999: Ceoorganopoulos Miuscher 1998)., 1999; Chiaberge Ghisellini 1999; Georganopoulos Marscher 1998). + Tf vy29ve. then fuum)Xtoot). and tag©fouu(ro). l0. the time lag of a flare at low frequency to its corresponding svuchrotron flare at high frequency is roughly the radiative cooling time of relativistic electrous at lower frequency (παν ct al.," If $\nu_1\gg \nu_2$ , then $t_{cool}(\nu_2)\gg t_{cool}(\nu_1)$, and $t_{lag} \approx t_{cool}(\nu_2)$, i.e., the time lag of a flare at low frequency to its corresponding synchrotron flare at high frequency is roughly the radiative cooling time of relativistic electrons at lower frequency (Urry et al." + 1997. 1999).," 1997, 1999)." + Tu the case that N-xavs are due to svuchrotron emission. svuchrotron cooling is the dominant cooling process of electrons. i.e. 1|D~1.," In the case that X-rays are due to synchrotron emission, synchrotron cooling is the dominant cooling process of electrons, i.e., $1+D\sim 1$." +" Thus. in the jet comoving frame where Bois in gauss. 45—ye/(10TIz) is in the observers frame. aud fj,4 1s iu units of second."," Thus, in the jet comoving frame where $B$ is in gauss, $\nu_{15}\equiv \nu_{2}/(10^{15} {\rm Hz})$ is in the observer's frame, and $t_{lag}$ is in units of second." +" The tine lags of optical (V. baud. observers frame) endssiou to svuchrotron ταν emissiou im the jet comoving frame are and the time lags of radio (1 Πε, observers frame) Cluission to A-ray enussion in the jet comoving frame are As the emitting plasiia. propagates downstream alone the jet. £,,;, and f,.7, result in offsets between the ceutroids of optical aud X-ray. knots. and between the ceutroids of radio and N-ray knots respectively, with the X-ray knot being earlier and closer to the core than radio aud optical knots. and optical knot than radio knot."," The time lags of optical (V band, observer's frame) emission to synchrotron X-ray emission in the jet comoving frame are and the time lags of radio (1 GHz, observer's frame) emission to X-ray emission in the jet comoving frame are As the emitting plasma propagates downstream along the jet, $t_{o/x}$ and $t_{r/x}$ result in offsets between the centroids of optical and X-ray knots, and between the centroids of radio and X-ray knots, respectively, with the X-ray knot being earlier and closer to the core than radio and optical knots, and optical knot than radio knot." + The radio/XN-rav or optical/X-ray knot pairs observed at the sale time are thus actually produced by two different populations of electrous. as pointed out by Siomiginowska et al. (," The radio/X-ray or optical/X-ray knot pairs observed at the same time are thus actually produced by two different populations of electrons, as pointed out by Siemiginowska et al. (" +2002).,2002). + IIowever. bright knots on large scales are sites continuously eeucratiug shocks and re-accelerating electrous. as mdicated by proper motion at knots of M87 jet (Diretta et al.," However, bright knots on large scales are sites continuously generating shocks and re-accelerating electrons, as indicated by proper motion at knots of M87 jet (Biretta et al." + 1995. 1999).," 1995, 1999)." + The maguetic field strength in bright knots aud the jet speed οὐ are roughly. coustaut (on time scale of vers)., The magnetic field strength in bright knots and the jet speed $v_j$ are roughly constant (on time scale of years). + Offsets produced. by a series of xo»pulatious of electrous at a bright knot are approxinatelv he same., Offsets produced by a series of populations of electrons at a bright knot are approximately the same. + That is to sav. the observed radio/N-rav iud optical/N-rav offsets at a bright knot still reflect the svuchrotron time lags. as i£ they were produced by a single »opulatiou of electrous at the same time.," That is to say, the observed radio/X-ray and optical/X-ray offsets at a bright knot still reflect the synchrotron time lags, as if they were produced by a single population of electrons at the same time." + Απο ej=0.12 (P—1.1. Colotti et al.," Assuming $v_j=0.42c$ $\Gamma=1.1$, Celotti et al." +" 2001) on arge scales. the observed offset between optical (V. band) and X-ray knots caused by £,;; simply is where Bo,=Ο10teauss)."," 2001) on large scales, the observed offset between optical (V band) and X-ray knots caused by $t_{o/x}$ simply is where $B_{-4}\equiv B/(10^{-4} {\rm gauss})$." +" Similarly, the observed offsets between radio (1 GIIz) aud X-ray knots caused by foods The typical magnetic strength in large scale jets is B~10! eauss (Iris Narwezvuski 2002)."," Similarly, the observed offsets between radio (1 GHz) and X-ray knots caused by $t_{r/x}$ is The typical magnetic strength in large scale jets is $B\sim 10^{-4}$ gauss (Harris Karwczynski 2002)." + According to equations (13) and (11). the de-projected opticalX-rav and radio/X-rav offsets of large scale jets are ~65.7 pe and ~SOkpe. respectively.," According to equations (13) and (14), the de-projected optical/X-ray and radio/X-ray offsets of large scale jets are $\sim 65.7$ pc and $\sim 89$ kpc, respectively." + However. jet expansion ou large scales mua shorten the svuchrotron time lags aud hence the offsets. especially radio/optical offsets.," However, jet expansion on large scales may shorten the synchrotron time lags and hence the offsets, especially radio/optical offsets." + Lower energev clectrons cool slower. hence at lower frequencies. jets have longer time to expand sideways and thus are wider.," Lower energy electrons cool slower, hence at lower frequencies, jets have longer time to expand sideways and thus are wider." + Multibaud imaging ofjets in Pictor A (Wilson et al., Multiband imaging of jets in Pictor A (Wilson et al. + 2001). Con A (ISraft et al.," 2001), Cen A (Kraft et al." + 2002). NIST (Marshall et al.," 2002), M87 (Marshall et al." + 2002: Sparks et al., 2002; Sparks et al. + 1996). and 3€ 273 (Thomson et al.," 1996), and 3C 273 (Thomson et al." + 1993) shows that radio jets are much wider. indicating that jet expansion on large scales is significant.," 1993) shows that radio jets are much wider, indicating that jet expansion on large scales is significant." + It is reasonable that 1e 1nagnetic field farther away from the jet axis is weaker., It is reasonable that the magnetic field farther away from the jet axis is weaker. + When the jet expands. some electrous move sidewavs ο region with weaker magnetic field. enüttiusg at lower yequencies earlier than they should.," When the jet expands, some electrons move sideways to region with weaker magnetic field, emitting at lower frequencies earlier than they should." + Iu other words. jet expansion reduces the time lags for these electrons. as well as the jet emission at some frequencies.," In other words, jet expansion reduces the time lags for these electrons, as well as the jet emission at some frequencies." + On large scales. he magnetic field is very weak. aud onlv at frequencies Close to Maa. is the jet expansion ucelieihle.," On large scales, the magnetic field is very weak, and only at frequencies close to $\nu_{max}$, is the jet expansion negligible." + At low requencies (2 sole cases may be as liehl as optical). the jet ma expand so ereatlv that the emission of the jet is uaiulv contributed bw clectrous iu the expanded region of the jet. aud cousequeutlv the ceutroids of emission at hese frequencies shift wpstream. ie. the offsets to high requeucy knots are shortened.," At low frequencies (in some cases may be as high as optical), the jet may expand so greatly that the emission of the jet is mainly contributed by electrons in the expanded region of the jet, and consequently the centroids of emission at these frequencies shift upstream, i.e., the offsets to high frequency knots are shortened." + Tn equation (9). it can be seen that at a given requency. svuchrotron time lags aud thus offsets between wo chussion frequencies are inversely proportional to the uaenetic field strength B.," In equation (9), it can be seen that at a given frequency, synchrotron time lags and thus offsets between two emission frequencies are inversely proportional to the magnetic field strength $B$." + Because B decreases aloug the jet. offsets due to svuchrotron time lags should increase along the jet provided that jet expansion is uuiniportaut.," Because $B$ decreases along the jet, offsets due to synchrotron time lags should increase along the jet provided that jet expansion is unimportant." + Iu MsS87. Con A. 3€ 273 and PINS 1127.115 nore than wo knots have been detected.," In M87, Cen A, 3C 273 and PKS $1127-145$ more than two knots have been detected." + The offsets in M87. Cen A and 3C 273 indeed ect larger along the jet. suggesting hat these offsets ave due to svuchrotron time lags.," The offsets in M87, Cen A and 3C 273 indeed get larger along the jet, suggesting that these offsets are due to synchrotron time lags." +" In MS8ST. roni knot A to Nuot D. the optical/N-ray offsets increase ron 0.09+0.02” to 0.294:0.06"". aud from the core to knot A. the optical/N-rav offsets are too sinall to be measured at current accuracy of 0.03” (NDarsball et al."," In M87, from knot A to Knot B, the optical/X-ray offsets increase from $0.09 \pm 0.02\arcsec$ to $0.29 \pm 0.06\arcsec$, and from the core to knot A, the optical/X-ray offsets are too small to be measured at current accuracy of $\arcsec$ (Marshall et al." + 2002)., 2002). +" The radio/optical offsets in ALS7 also ect larger along the jet. from O.2140.01"" at knot IIST-1 to ~0.:3250.01"" at knot D. Iu Ceu A. the radio/N-rav offsets increase frou <0.5"" at AL/AXI knots aud ~2.5"" at A?/AN2 knots to ~5"" at knot D (Ixraft et al."," The radio/optical offsets in M87 also get larger along the jet, from $\sim0.24\pm 0.01\arcsec$ at knot HST-1 to $\sim0.3\pm 0.01\arcsec$ at knot D. In Cen A, the radio/X-ray offsets increase from $<0.5\arcsec$ at A1/AX1 knots and $\sim2.5\arcsec$ at A2/AX2 knots to $\sim5\arcsec$ at knot B (Kraft et al." + 2002)., 2002). +" In 3€ 273. the optical/X-rav offsets ineroase fron ~0.22"" at knot A to ~0.5"" at knot B (see Fig."," In 3C 273, the optical/X-ray offsets increase from $\sim0.22\arcsec$ at knot A to $\sim0.5\arcsec$ at knot B (see Fig." + 3 in Marshall et al., 3 in Marshall et al. + 2001)., 2001). + The radio/X- offsets im PINS 1127.115 and the radio/optical offsets bevoud knot D in M87 decrease along the jet. but this may be due to strong expansion of the jet.," The radio/X-ray offsets in PKS $1127-145$ and the radio/optical offsets beyond knot D in M87 decrease along the jet, but this may be due to strong expansion of the jet." + Iu fact. from pe aud subpe scales to large scales. offsets also increase along the jet.," In fact, from pc and subpc scales to large scales, offsets also increase along the jet." + The typical time lag of optical flares to N-ravs flares is of dav scale. such as in PIS 9155301 (Uy et al.," The typical time lag of optical flares to X-rays flares is of day scale, such as in PKS $2155-304$ (Urry et al." + 1997) and Mrk 121 (Buckley et al., 1997) and Mrk 421 (Buckley et al. + 1996)., 1996). + The corresponding optical/N-ray offsets ou pe and subpe scales are of the order of lightday. ic. thousaudtlis of pc.," The corresponding optical/X-ray offsets on pc and subpc scales are of the order of lightday, i.e. thousandths of pc." + On large scales. the observed opticalταν offsets iu large scale jets are of the order of tens (M87. assuniug 30°) to hundreds (060 273 aud 3€ 66B) of pe.," On large scales, the observed optical/X-ray offsets in large scale jets are of the order of tens (M87, assuming $\theta\sim 30\arcdeg$ ) to hundreds (3C 273 and 3C 66B) of pc." + The typicaltime lag of radio flares to optical flares is of the scale of inoutlis (e.e¢.. Tornikoski et al.," The typicaltime lag of radio flares to optical flares is of the scale of months (e.g., Tornikoski et al." + 1991). corresponding to radio/optical aud racdio/X-ray offsets of the order of 0.1 pe on pe aud subpe scales. while the observed," 1994), corresponding to radio/optical and radio/X-ray offsets of the order of 0.1 pc on pc and subpc scales, while the observed" + 1991). corresponding to radio/optical aud racdio/X-ray offsets of the order of 0.1 pe on pe aud subpe scales. while the observed.," 1994), corresponding to radio/optical and radio/X-ray offsets of the order of 0.1 pc on pc and subpc scales, while the observed" +uctallicities in the rauge of |Fe/II| 2|-3. 0].,"metallicities in the range of [Fe/H] =[-3, 0]." + Of these stars. approximately are in the NW quadrant which is our area of interest.," Of these stars, approximately are in the NW quadrant which is our area of interest." + The ealactic stars that happen o fall iuto the range of colors and magnitudes that are appropriate for MOI red eiauts create a slowly varvine oreeround which will be removed in the analysis of the stream density., The galactic stars that happen to fall into the range of colors and magnitudes that are appropriate for M31 red giants create a slowly varying foreground which will be removed in the analysis of the stream density. + Note that a shorter distance would make he interred Iuninosities simaller aud cause the isochrone natch to systematically vield larger metallicities., Note that a shorter distance would make the inferred luminosities smaller and cause the isochrone match to systematically yield larger metallicities. + To a first approximation this simply offsets the metallicity range we analyze and docs not otherwise atfect the analysis., To a first approximation this simply offsets the metallicity range we analyze and does not otherwise affect the analysis. + Star streams are identified as spatially cohereut over-densities. over a linited rauge of metallicities.," Star streams are identified as spatially coherent over-densities, over a limited range of metallicities." + To create the map we project the RGB stars onto a tangent plane. for which we chose a center at RA=Oh 2112 and Dec = wwhich is near the ceuter of the NW stream.," To create the map we project the RGB stars onto a tangent plane, for which we chose a center at RA=0h 24 m and Dec = which is near the center of the NW stream." + This center ds offset from ΑΟ center. to make the co-ordinates reasonably rectangular in the region of the stream analysis.," This center is offset from M31's center, to make the co-ordinates reasonably rectangular in the region of the stream analysis." + Without the ROB star selection no streams are readily visible iu the halo., Without the RGB star selection no streams are readily visible in the halo. + Figure 1. plots the sky distribution (in the tangent plane co-ordinates) of the [Fe/TI|2|-0.6. -2.1] stars which we find below comprise the stream.," Figure \ref{fig_M31} plots the sky distribution (in the tangent plane co-ordinates) of the [Fe/H]=[-0.6, -2.4] stars which we find below comprise the stream." + The same feature is clear m auv map mace with the low metallicity star sample., The same feature is clear in any map made with the low metallicity star sample. + The low aetallicity map. Figuree νι has a umuber of ecolctric sampling variations that must be coutrolled ο 1uake a reliable measurement of local stream ceusity.," The low metallicity map, Figure \ref{fig_M31}, has a number of geometric sampling variations that must be controlled to make a reliable measurement of local stream density." +" The simplest problem is that the CCD array has gaps )etwoeen the individual chips of 80"" iiu the horizoutal direction aud ~13” iu the vertical direction.", The simplest problem is that the CCD array has gaps between the individual chips of $\sim$ in the horizontal direction and $\sim$ in the vertical direction. + The images were dithered ο cover the smaller gaps. but the resulting depth is shallower than the surrounciug area aud no simple count correction procedure that did not inerease the uoise was ound.," The images were dithered to cover the smaller gaps, but the resulting depth is shallower than the surrounding area and no simple count correction procedure that did not increase the noise was found." + These gaps are identified in the camera CCD co-ordinate svstem and masked out of the map., These gaps are identified in the camera CCD co-ordinate system and masked out of the map. + A more complicated geometric bias results from the catalog bemg ilt up from stacked local mages in which the stellar photometry was done to create a local catalog., A more complicated geometric bias results from the catalog being built up from stacked local images in which the stellar photometry was done to create a local catalog. + These catalogs were then matched in the areas of overlap to chininate duplicates and produce a global catalog., These catalogs were then matched in the areas of overlap to eliminate duplicates and produce a global catalog. + The nuages are vienetted towards the edee of the array so the image depth drops at the field edge. however he vignetting is partially compensated in the catalog natching at the nuage edges.," The images are vignetted towards the edge of the array so the image depth drops at the field edge, however the vignetting is partially compensated in the catalog matching at the image edges." + That is. if the surface density is corrected with the inverse of the mean star density over the array then the map has an excess at he edges.," That is, if the surface density is corrected with the inverse of the mean star density over the array then the map has an excess at the edges." + It is straightforward to devise a correction which flattens the average surface density., It is straightforward to devise a correction which flattens the average surface density. + The correction works well over a range of about 1n surface density. ot undesirably amplifies the noise if applied at lower conipleteness levels.," The correction works well over a range of about in surface density, but undesirably amplifies the noise if applied at lower completeness levels." + We apply a cut to the image where he local surface density in CCD nuage co-ordinates drops below of the average., We apply a cut to the image where the local surface density in CCD image co-ordinates drops below of the average. + The resulting masked out area is about of the image. a fraction which is rot very sensitive to the precise value of the cut.," The resulting masked out area is about of the image, a fraction which is not very sensitive to the precise value of the cut." + The nasked out regious are not included iu the analysis aud are readily visible in Figure 1.., The masked out regions are not included in the analysis and are readily visible in Figure \ref{fig_M31}. + Field to field photometric depth variatious will introduce artificial variations in the deusitv of stars in the map., Field to field photometric depth variations will introduce artificial variations in the density of stars in the map. + The depth of an image for a fixed exposure time depends on the sky transparency. brightucss and tage quality at the time of observation.," The depth of an image for a fixed exposure time depends on the sky transparency, brightness and image quality at the time of observation." + Below he completeness limit there is substantial field to fell variation that leads to a random checkerboard appearance. but the pattern esscutially disappears above he completeness limit.," Below the completeness limit there is substantial field to field variation that leads to a random checkerboard appearance, but the pattern essentially disappears above the completeness limit." + Our measurement of the stream density subtracts the local mean backeround to reduce he problems of local depth variations., Our measurement of the stream density subtracts the local mean background to reduce the problems of local depth variations. + A second. more difficult. form of depth variations is fell to field variatious in the zero points of the two filter bands.," A second, more difficult, form of depth variations is field to field variations in the zero points of the two filter bands." + We can set a Init ou these by usine the double measurements of star brightuess in the overlap reeions., We can set a limit on these by using the double measurements of star brightness in the overlap regions. + For bright stars where the photon noise is small he standard deviation is 0.037 mae in g baud aud 0.052 nag in / band. with no disceruible pattern over the field.," For bright stars where the photon noise is small the standard deviation is 0.037 mag in $g$ band and 0.052 mag in $i$ band, with no discernible pattern over the field." + The star catalog in our selected rauge of color aud uaeguitude has nearly coustaut nuubers with increasing depth near the completeness limit. so a chanee iu zero point leads to about a change in uuubers.," The star catalog in our selected range of color and magnitude has nearly constant numbers with increasing depth near the completeness limit, so a change in zero point leads to about a change in numbers." + The should be considered au upper ΠΕ since the edges of the images are vienetted aud have significant local calibration variations., The should be considered an upper limit since the edges of the images are vignetted and have significant local calibration variations. + Across cach \MAleegaprime field there are a iuvriad of camera ordinate dependent photometric offsets (Beenaultetal.2009) which the current calibration does not fully take iuto account., Across each Megaprime field there are a myriad of camera co-ordinate dependent photometric offsets \citep{Regnault:09} which the current calibration does not fully take into account. + These offsets lead to a change im zero point witlin each one degree pointing that is an additional ~2- variation., These offsets lead to a change in zero point within each one degree pointing that is an additional $\sim$ variation. + The masking procedure removes tle worst, The masking procedure removes the worst +for the thick shell (i.e. relativistic) case.,for the thick shell (i.e. relativistic) case. + Here. the GRB duration T is given bv the shel width Aj/e. according to the internal shock model.," Here, the GRB duration T is given by the shell width $\Delta_0/c$, according to the internal shock model." +" As derived by Sari Piran (19994) and Kobayashi (2000). the twpical svnchrotron frequency and peak fInx evolve asMyx1°. FosseX073/2 a Em Lus and puXE-51/35OPP Eusox-dAl/357879 at Eo-Lus forj the thin-shel: case: py, —constant. Foor,oxIU? ab Lm Teand pusxdDIE uuuxd.PI abdo T. [or the thick-shell case."," As derived by Sari Piran (1999a) and Kobayashi (2000), the typical synchrotron frequency and peak flux evolve as$\nu_{m,r}\propto t^6$ , $F_{\nu,max,r}\propto +t^{3/2}$ at $tt_\times$ , for the thin-shell case; $\nu_{m,r}=$ constant, $F_{\nu,max,r}\propto t^{1/2}$ at $tT$, for the thick-shell case." +" Before the crossing time. the cooling frequency evolve as pi,x17 for the thin-shell case and 1,x/+ for the thick-shell case: alter the crossing time. 7, (urns into the cutoff frequency due to svuchrotvon cooling of electrons without acceleration. and it evolves the same as η, lor both cases."," Before the crossing time, the cooling frequency evolve as $\nu_{c,r}\propto t^{-2}$ for the thin-shell case and $\nu_{c,r}\propto t^{-1}$ for the thick-shell case; after the crossing time, $\nu_{c,r}$ turns into the cutoff frequency due to synchrotron cooling of electrons without acceleration, and it evolves the same as $\nu_{m,r}$ for both cases." +" According to the emission properties shown above. we derive light curves of the emission from the forward- and reverse-shocked regions and thus can constrain the model parameters by fitting the early afterglow of GRD 0505254. From the emission features described by equations (1). (2). (5). (6). (8). aud (9). we see that the optical band is commonly between the (vpical Lequency 74, and cooling Irequency D. m both forward and reverse shocks around the crossing Gime."," According to the emission properties shown above, we derive light curves of the emission from the forward- and reverse-shocked regions and thus can constrain the model parameters by fitting the early afterglow of GRB 050525A. From the emission features described by equations (1), (2), (5), (6), (8), and (9), we see that the optical band is commonly between the typical frequency $\nu_m$ and cooling frequency $\nu_c$ in both forward and reverse shocks around the crossing time." +" Thus. the optical light curve of GRB 050525AÀ can be fitted by the superposition of two broken power-laws: (1) forward. shock: Ενα{ο1/2 lor. lyUypt (2)- reverse shock: ΕνXE,B(2p?1/2 for,. ly<» ly. and PF»,xf,( for fy>fy."," Thus, the optical light curve of GRB 050525A can be fitted by the superposition of two broken power-laws: (1) forward shock: $F_{\nu,f}\propto t_d^{1/2}$ for $t_d < t_{m,f}$, and $F_{\nu,f}\propto +t_d^{3(1-p)/4}$ for $ t_d > t_{m,f}$; (2) reverse shock: $F_{\nu,r}\propto +t_d^{3(2p-1)/2}$ for $t_d < t_\times$ , and $F_{\nu,r}\propto t_d^{(-27p+7)/35}$ for $ t_d > t_\times$." + These scaling laws are valid for the thin-shell case. in which the reverse shock decelerates the shell insignificantly (IXobavashi 2000)..," These scaling laws are valid for the thin-shell case, in which the reverse shock decelerates the shell insignificantly \citep{k00}. ." + InFigure 1 wepresent the R-bandalterglow light curve of GRB 0505254 between Alay 25.0062 UT (Af220.004 das) and May 25.5333 UT CM0.531 clavs).," InFigure \ref{fig1} wepresent the R-bandafterglow light curve of GRB 050525A between May 25.0062 UT $\Delta t\approx 0.004\,\,{\rm days}$ ) and May 25.5333 UT $\Delta t\approx 0.531\,\,{\rm days}$ )." + We findthat.," We findthat," +telescope itself. thereby avoiding Llexing of the spectrograph unit as the telescope follows the target field across the sky.,"telescope itself, thereby avoiding flexing of the spectrograph unit as the telescope follows the target field across the sky." + The object exposures were combined to form a single object image. which was then processed by theIRAF taskDOUYDRA.," The object exposures were combined to form a single object image, which was then processed by the task." + This task uses the combined sky-lat image to identify the position of the spectra in the object. image. although it does require manual intervention at times. as the spectra on the object image are 13 pixels wide with only à ~2 pixel separation.," This task uses the combined sky-flat image to identify the position of the spectra in the object image, although it does require manual intervention at times, as the spectra on the object image are 1–3 pixels wide with only a $\sim2$ pixel separation." + The relatively even illumination across the twilight sky image ensures that the spectra are correctly traced from the red wavelengths to the blue., The relatively even illumination across the twilight sky image ensures that the spectra are correctly traced from the red wavelengths to the blue. + Third-order polynomials are quite adequate to fit the trace of the spectra. with tvpical rms errors being 0.1.0.2 pixels.," Third-order polynomials are quite adequate to fit the trace of the spectra, with typical rms errors being 0.1–0.2 pixels." + The combined [at-feld frames were then απο o calculate the relative. fibre. throughput. for which compensations are made in the extracted. spectra.," The combined flat-field frames were then used to calculate the relative fibre throughput, for which compensations are made in the extracted spectra." + Fibre hroughput tends to degrade with time as dust and. erime accumulate on the prisms., Fibre throughput tends to degrade with time as dust and grime accumulate on the prisms. + As the system ages. the elue »'ween the prism and fibre becomes more opaque. ancl accidental kinks in the fibres take their toll.," As the system ages, the glue between the prism and fibre becomes more opaque and accidental kinks in the fibres take their toll." + Alisalignment in the mounting of the prism on the fibre buttons due o maintenance repairs or knocks in positioning. ancl poor seating of the fibre button the plate during configuration. also reduce a fibres throughput. as do telescope vignetting and the variation in transmission of a [fibre from. the manufacturing process.," Misalignment in the mounting of the prism on the fibre buttons due to maintenance repairs or knocks in positioning, and poor seating of the fibre button the plate during configuration, also reduce a fibre's throughput, as do telescope vignetting and the variation in transmission of a fibre from the manufacturing process." + The wavelength calibration is based. on identifying 225 lines in the are spectra and fitting these with thirc or fourth-order polvnomials to obtain the wavelengthpixel relation., The wavelength calibration is based on identifying 20--25 lines in the arc spectra and fitting these with third- or fourth-order polynomials to obtain the wavelength–pixel relation. + “Pvpical rms errors in the fits are about 0.1 pixels. or about0.," Typical rms errors in the fits are about 0.1 pixels, or about." +5A.. The individual sky spectra are then inspected. and any spectrum contaminated. by light bleeding from. bright objects in neighbouring fibres is removed before the sky spectra are. combined.," The individual sky spectra are then inspected, and any spectrum contaminated by light bleeding from bright objects in neighbouring fibres is removed before the sky spectra are combined." + Although potentially a serious problem. contamination from neighbouring fibres is quite rare. and tends to manifest as a mocdification of the spectral continuum of both ojects and sky.," Although potentially a serious problem, contamination from neighbouring fibres is quite rare, and tends to manifest as a modification of the spectral continuum of both objects and sky." + This does not alfect 10 position of spectral lines. but onlv the strength of the —ines relative to the continuum: for measuring redshifts it is therefore not a particular problem except where the sky Es»eetrum. and hence sky subtraction. is allected.," This does not affect the position of spectral lines, but only the strength of the lines relative to the continuum; for measuring redshifts it is therefore not a particular problem except where the sky spectrum, and hence sky subtraction, is affected." + Phere is often considerable variation in the sky. spectra due to the Iarge field of view ofthe Schmidt telescope. which results in the sky fibres sampling parts of the skv up to aapart.," There is often considerable variation in the sky spectra due to the large field of view of the Schmidt telescope, which results in the sky fibres sampling parts of the sky up to apart." + Most. of the variations are in the strengths of the garong emission features at5577... aand63634... while the continuum tends to remain unadIectecd.," Most of the variations are in the strengths of the strong emission features at, and, while the continuum tends to remain unaffected." + This results in good continuum subtraction. but poor subtraction of the stronglv-varving sky lines. which can lead to cillieulties in identifving spectral features near these lines.," This results in good continuum subtraction, but poor subtraction of the strongly-varying sky lines, which can lead to difficulties in identifying spectral features near these lines." + Once the spectra have been extracted. sky-subtracted. and wavelength-calibrated. radial velocity (recishift) measurements can be made.," Once the spectra have been extracted, sky-subtracted, and wavelength-calibrated, radial velocity (redshift) measurements can be made." + The CCL on the FLAIR system prior to. 1995 had a rather poor response in the rluc. making it. dillieult to identify the Ca LI and Ix lines and the €? band.," The CCD on the FLAIR system prior to 1995 had a rather poor response in the blue, making it difficult to identify the Ca H and K lines and the G band." + While the new thinned CCD is twice as responsive in the blue. these lines can still be quite dilficult o detect.," While the new thinned CCD is twice as responsive in the blue, these lines can still be quite difficult to detect." + Other lines used in the process of measuring the redshift were H3 at486LA.. at aancd5OOTA.. the Mg b. triplet. around5175... the be eatures at aand5270A.. Na D atSSOPA.. Ho atG5O3A.. atGOSLA.. and the doublet at aand631A.," Other lines used in the process of measuring the redshift were $\beta$ at, at and, the Mg b triplet around, the Fe features at and, Na D at, $\alpha$ at, at, and the doublet at and." +.. Using the taskHVIDLINES. gaussians were fitted to each identified feature in the spectrum. to. determine the position of the line and. hence its shift.," Using the task, gaussians were fitted to each identified feature in the spectrum to determine the position of the line and hence its shift." + Ehe line shifts are then averaged to produce a final racial velocity measurement. while the standard deviation of the line shifts eives the internal racial velocity error.," The line shifts are then averaged to produce a final radial velocity measurement, while the standard deviation of the line shifts gives the internal radial velocity error." + These errors were typically less than LOO ((see below)., These errors were typically less than 100 (see below). + Additional redshifts for the ΕΙΛΑΡΙ survev catalogue were obtained from the literature via the NASA/IPAC Extragalactic Database (NED) and the ZCAT catalogue (Lluchra19001 901 redshifts were obtained from εςΑΕ. and 354 were obtained [rom NED.," Additional redshifts for the FLASH survey catalogue were obtained from the literature via the NASA/IPAC Extragalactic Database (NED) and the ZCAT catalogue \cite{zcat}; 901 redshifts were obtained from ZCAT, and 354 were obtained from NED." + The final estimate of the galaxv's redshift. is. the variance-weightec mean of the measured. recshilts [rom FLAIR (including repeat measurements for some objects) and the literature., The final estimate of the galaxy's redshift is the variance-weighted mean of the measured redshifts from FLAIR (including repeat measurements for some objects) and the literature. + X. further. subjective. weighting was applied to the literature-derived: redshifts. rellecting our degree. of belief. in the estimated. redshift) errors in the literaturein short. NED and ZCXT redshifts were given half the weight of the FLALR redshifts.," A further, subjective, weighting was applied to the literature-derived redshifts, reflecting our degree of belief in the estimated redshift errors in the literature—in short, NED and ZCAT redshifts were given half the weight of the FLAIR redshifts." + One can check the precision of the FLALR. recshilts by comparing the multiple radial velocity measurements for those galaxies that have repeat observations., One can check the precision of the FLAIR redshifts by comparing the multiple radial velocity measurements for those galaxies that have repeat observations. + The recshift differences for the 164. galaxies with multiple redshift measurements are plotted in the top panel of Figure 2.., The redshift differences for the 164 galaxies with multiple redshift measurements are plotted in the top panel of Figure \ref{fig:deltav}. + Phe mean olfset is O4ts: the standard deviation in the dillerence is 134 indicating an rms error in the redshifts of 95 ((assumingὃν equipartition of errors).," The mean offset is $-$ 0.4; the standard deviation in the difference is 134, indicating an rms error in the redshifts of 95 (assuming equipartition of errors)." + Although we have been careful in. estimating the redshift errors. there mav be cllects (systematic or otherwise) that have not been accounted for.," Although we have been careful in estimating the redshift errors, there may be effects (systematic or otherwise) that have not been accounted for." + To detect such ellects and determine the reliability ofthe estimated errors. one can compare the variance-weighted estimated errors. 7. with the rms errors. ejt for the 164 objects with repeat measurements.," To detect such effects and determine the reliability ofthe estimated errors, one can compare the variance-weighted estimated errors, $\overline{\sigma}$, with the rms errors, $\sigma_{\rm RMS}$ , for the 164 objects with repeat measurements." + , +The spectropolarimeter SOT/SP (Litesctal.2001:Tsunetaetal.2008) aboard the Japanese satellite (Ixosueietal.2007) combines (17332 angular resolution and 10? polarimetric scusitivity to perform seeiue-free full Stokes 1icasuremenuts of the polarized light emereinme frou the maguetized solar plotosphere.," The spectropolarimeter SOT/SP \citep[][]{Lit01,Tsu08} + aboard the Japanese satellite \citep{Kos07} + combines 32 angular resolution and $10^{-3}$ polarimetric sensitivity to perform seeing-free full Stokes measurements of the polarized light emerging from the magnetized solar photosphere." + Since its launch in 2006. the SOT/SP iustrumeut has allowed the solar commmuity to investigate the solar surface miaegnetiqu under unprecedented stable conditions.," Since its launch in $2006$, the SOT/SP instrument has allowed the solar community to investigate the solar surface magnetism under unprecedented stable conditions." + SOT/SP observations have been larecly exploited to investigate the quiet Sun magnetis., SOT/SP observations have been largely exploited to investigate the quiet Sun magnetism. + Its elobal properties Lave been described by OrozcoSuárezetal. (2007a).. Litesetal. 2008)... ÀsensioRamos (2009).. and Jinetal.(2009): while detailed analyses of its local properties. in relation with the temporal evolution due to the interaction with plasma motions. have been carried out by Ceutenoetal.(2007).. OrozcoSuárezetal. (2008)... Nagataetal. (2008).. Shimizuetal. (2008).. Fischeretal.(2009) and Zhangetal.(2009).," Its global properties have been described by \citet{OroS07}, \citet{Lit08}, \citet{AseR09}, and \citet{Jin09}; while detailed analyses of its local properties, in relation with the temporal evolution due to the interaction with plasma motions, have been carried out by \citet{Cen07}, \citet{OroS08}, \citet{Nag08}, \citet{Shi08}, \citet{Fis09} and \citet{Zha09}." +. The above studies wave been perforned exploiting inversion techniques for interpretation of the observed Stokes profiles., The above studies have been performed exploiting inversion techniques for interpretation of the observed Stokes profiles. + Iu mos of the cases. the hypothesis of ME atinosphere has heen adopted to infer the properties of the magnetic field vector iu vosolution clemeuts.," In most of the cases, the hypothesis of ME atmosphere has been adopted to infer the properties of the magnetic field vector in resolution elements." + It assumes that the polarization is produced i an atimosphere where the magnetic field vector is constant., It assumes that the polarization is produced in an atmosphere where the magnetic field vector is constant. + This approximation ix a good compromise to iuterpret aree data sets iu reasonable times. as the ones obtained bvHINODE.," This approximation is a good compromise to interpret large data sets in reasonable times, as the ones obtained by." +. However. it is nuportaut to realizo that. aneular resolution is still too low to perform spectropoludimetrie observations of iiagnetic structures that can be regarded as uniform (c.g..SAuchezAhueida2001:Stenflo 2009).," However, it is important to realize that angular resolution is still too low to perform spectropolarimetric observations of magnetic structures that can be regarded as uniform \citep[e.g.,][]{SanA04b,Ste09}." +.. As we cunphasize in 3.1. asyhunetric Stokes profiles in quiet. Sun observatious are he rule rather thau he exception. which nuples the presence of unresolved structure (SánchezAlmeidaetal.1996.andreferences therein).," As we emphasize in \ref{resexam}, asymmetric Stokes profiles in quiet Sun observations are the rule rather than the exception, which implies the presence of unresolved structure \citep[][and references therein]{SanA96}." + The ME analyses cannot reproduce them. overlooking: the expected coexistence of several magnetic componcuts iu a single resolution clemeut +.," The ME analyses cannot reproduce them, overlooking the expected coexistence of several magnetic components in a single resolution element ." + SáuchezAhuecidaetal.(1996) put forward iveuimoeuts for describing the atimosphere where the polarization is formed as a MIero-Structured. Alaenetic Atinosphlere (MISMA). Le. an atinosphere in which maguetic fields have structure simaller than the photon moean-fee path at the solar photosphere C5100 ln).," \citet{SanA96} put forward arguments for describing the atmosphere where the polarization is formed as a MIcro-Structured Magnetic Atmosphere (MISMA), i.e., an atmosphere in which magnetic fields have structure smaller than the photon mean-free path at the solar photosphere $\la100$ km)." + I can be regarded as a Πήλιο case of structures of all. sizes (e.g.LaudiDeelIunoceuti199 la.b)..," It can be regarded as a limiting case of structures of all sizes \citep[e.g.][]{lan94,car07}." + The MISMA. livpotliesis siuplif&es the radiative trausfer. vet it provides realistic asvuuuetric Stokes profiles.," The MISMA hypothesis simplifies the radiative transfer, yet it provides realistic asymmetric Stokes profiles." + Uuder this hvpothesis. the polarization from a siugle pixel in spectropolarimnetric observations is equivalent to that produced by the average atinosphere.," Under this hypothesis, the polarization from a single pixel in spectropolarimetric observations is equivalent to that produced by the average atmosphere." +" If one considers several conrponeuts with diverse physical properties (νο, thoermodyvuauies. plaza iuotions. and magnetic fields). the resulting spectrum is not a linear combination of Stokes profiles clnereine from cach magnetic component."," If one considers several components with diverse physical properties (i.e., thermodynamics, plasma motions, and magnetic fields), the resulting spectrum is not a linear combination of Stokes profiles emerging from each magnetic component." + Rather. the superposition is non-linear. giviug rise to asviunietries with the properties observed in the Sun (e$. with spectral lines that produce net circular polarization).," Rather, the superposition is non-linear, giving rise to asymmetries with the properties observed in the Sun (e.g., with spectral lines that produce net circular polarization)." + SáuchezAlmeida&Lites(2000) showed Low such a colplex scenario can be properly adapted using a three coniponent model MISMA. whose implementation in an inversion code (SánchezAlueida1997) allowed them to reproduce the full —variety of profile asviuuetries emiereius from the quiet Sun when observed with the iustrunentation available at the time.," \citet{SanALit00} showed how such a complex scenario can be properly adapted using a three component model MISMA, whose implementation in an inversion code \citep{SanA97} allowed them to reproduce the full variety of profile asymmetries emerging from the quiet Sun when observed with the instrumentation available at the time." +" Stokes profiles iu IN auc network observations performed with exhibit very important asvuuuetries which encouraged us to attempt the same MISMA, analysis on these data.", Stokes profiles in IN and network observations performed with exhibit very important asymmetries which encouraged us to attempt the same MISMA analysis on these data. + We present an orderly inversion of Stokes profiles observed bv — SOT/SP iu IN and network regions., We present an orderly inversion of Stokes profiles observed by SOT/SP in IN and network regions. + It isperformed using the inversion code in Sanchez ).. that has been recently emiplowed to," It isperformed using the inversion code in \citet{SanA97}, , that has been recently employed to" +" T, = + +1- u— +P"," T_k = + +, = +." +"ron,Lia’: This is three equations for the three μι."," This is three equations for the three $\delta x_{k,1}$." +" The rest of the ory, then follow [rom equations (À)) and (AL3))."," The rest of the $\delta x_{k,n}$ then follow from equations \ref{eq:iteq}) ) and \ref{fixcoeffs}) )." +" For .N, particles and n, (me steps the computation time to get the 0:7;;, for all particles scales as NONwe"," For $N_p$ particles and $n_x$ time steps the computation time to get the $\delta x_{i,k,n}$ for all particles scales as $N_p^2N_x$." + In previous applications of the numerical action method (P9 ancl references therein) equation (À)) is solved by matrix inversion. an operation that scales as Αρη for relaxation of one particle orbit at a time.," In previous applications of the numerical action method (P9 and references therein) equation \ref{eq:tosolve}) ) is solved by matrix inversion, an operation that scales as $N_p^2n_x^3$ for relaxation of one particle orbit at a time." +" The NNAAI operation seales in the same war as a conventional numerical integration. but (he numerical prefactor is considerably larger because it takes many iterations of the coordinate shifts 0:7;4,, lo drive the 5,5, to zero."," The NNAM operation scales in the same way as a conventional numerical integration, but the numerical prefactor is considerably larger because it takes many iterations of the coordinate shifts $\delta x_{i,k,n}$ to drive the $S_{i,k,n}$ to zero." + NNAAI certainly will not replace conventional forward numerical integration for simulations ol the erowth of cosmic structure., NNAM certainly will not replace conventional forward numerical integration for simulations of the growth of cosmic structure. + But the application presented here shows why NNAÀM will be useful for analvses of the dvnamies of the nearby galaxies., But the application presented here shows why NNAM will be useful for analyses of the dynamics of the nearby galaxies. +"The advent of large galaxy surveys like the Sloan Digital Sky Survey (SDSS) in which photometry (and therefore colours) are readily. available for millions of objects has Lead to the common use of optical colours to define. ""earlv andlate"" tvpe galaxy samples (e.g.Simonetal.2009:SalimbeniCooray 2005).","The advent of large galaxy surveys like the Sloan Digital Sky Survey (SDSS) in which photometry (and therefore colours) are readily available for millions of objects has lead to the common use of optical colours to define “early"" and“late"" type galaxy samples \citep[e.g.][]{S09,S08,LP07,C07,B06,C05}." +. This method is particularly favoured since obtaining morphologics for large numbers of galaxies has until recently been impossible., This method is particularly favoured since obtaining morphologies for large numbers of galaxies has until recently been impossible. + Εις simplification is justified since it has been shown many times that the majority of galaxies follow. a strict colour-morphologv. relation., This simplification is justified since it has been shown many times that the majority of galaxies follow a strict colour-morphology relation. + For example Mignolietal.(2009) arguedl that of galaxies to ze] are either red. bulge-dominatedl galaxies or blue. disk dominated galaxies: while Consclice(2006) showed a similar result for 22000 low recshilt galaxies (both using automated methods for morphological classification).," For example \citet{M09} argued that of galaxies to $\sim$ 1 are either red, bulge-dominated galaxies or blue, disk dominated galaxies; while \citet{C06} showed a similar result for 22000 low redshift galaxies (both using automated methods for morphological classification)." + llowever the clear correlation between colour. and morpholoew is surprising. given that the colours of galaxies are determined. primarily by their. stellar content (and jorefore their recent star formation history. mostly within rw last. Civr) while the morphology is primarily clriven ov the dynamical history.," However the clear correlation between colour and morphology is surprising, given that the colours of galaxies are determined primarily by their stellar content (and therefore their recent star formation history, mostly within the last Gyr) while the morphology is primarily driven by the dynamical history." + The clear link between colour nd. morphology then gives ao strong indication that )0 timescales ancl processes which drive morphological ransformation and the cessation of star formation are garonely related - at least in most cases., The clear link between colour and morphology then gives a strong indication that the timescales and processes which drive morphological transformation and the cessation of star formation are strongly related - at least in most cases. + In this paper rowever. we consider a class of object (the red spirals) where 1e link described above appears to be broken.," In this paper however, we consider a class of object (the red spirals) where the link described above appears to be broken." + Since. the morpholosv-densitv relation was first quantified (Dressler1980) many mechanisms have been ooposed for the transformation of blue. star forming. clisk ealaxies in low density regions of the universe. to red. massive. earlv (vpe galaxies in clusters.," Since the morphology-density relation was first quantified \citep{D80} many mechanisms have been proposed for the transformation of blue, star forming, disk galaxies in low density regions of the universe, to red, passive, early type galaxies in clusters." + X recent review of many of the proposed mechanisms. and the evidence supporting them. can be found in Boselli&CGavazzi(2006).," A recent review of many of the proposed mechanisms, and the evidence supporting them, can be found in \citet{BG06}." +. Clearly. two things must happen for a star forming blue spiral galaxy to turn into a passive red early type., Clearly two things must happen for a star forming blue spiral galaxy to turn into a passive red early type. + First. star formation must cease (which can indirectlv alter the morphology by causing spiral armis ancl the disk in general to fade. possibly producing an SO or lenticular from a spiral). ancl secondly. in order to produce a ellipical the same. or a cilferent process must also dynamically alter the stellar kinematics of the galaxy.," First, star formation must cease (which can indirectly alter the morphology by causing spiral arms and the disk in general to fade, possibly producing an S0 or lenticular from a spiral), and secondly, in order to produce a elliptical the same, or a different process must also dynamically alter the stellar kinematics of the galaxy." + The presence of an unusually red or passive(ie., The presence of an unusually red or passive. +.. non- forming) population of spiral galaxies in clusters. of ealaxies was first noticed by vandenBereh(1010) in the Virgo cluster., non-star forming) population of spiral galaxies in clusters of galaxies was first noticed by \citet{V76} in the Virgo cluster. + Later studies of distant. cluster galaxics in (LIST) imaging also. revealed a significant number of so-called: “passive” spiral galaxies with a lack of on-going star formation (Couch et al., Later studies of distant cluster galaxies in (HST) imaging also revealed a significant number of so-called “passive” spiral galaxies with a lack of on-going star formation (Couch et al. + 1998: Dressler et al., 1998; Dressler et al. + 1999: Pogeianti ct al., 1999; Poggianti et al. + 1999)., 1999). + Passive late-type ealaxies were identified at lower redshifts in the outskirts of SDSS clusters by Gotoctal.(2003).. using concentration as a proxy for morphology.," Passive late-type galaxies were identified at lower redshifts in the outskirts of SDSS clusters by \citet{G03}, using concentration as a proxy for morphology." + Passive spirals in a cluster at zoO.4 were studied by Moranetal.(2006) who found star formation histories from GALEN observations consistent with the shutting down of star formation from strangulation (as described by Bekkictal. 2002))., Passive spirals in a cluster at $z\sim0.4$ were studied by \citet{M06} who found star formation histories from GALEX observations consistent with the shutting down of star formation from strangulation (as described by \citealt{Bekki02}) ). +" Passive spirals have also been revealed in a cluster at z—0.1 in the STAGES survey, using LIS’ morphologies Wolfetal.(2009)... rest [rame NUV-optical SEDs (Wolfetal.2)05) and 245m data [rom Spitzer (Callazzietal.2009)."," Passive spirals have also been revealed in a cluster at $z\sim 0.1$ in the STAGES survey, using HST morphologies \cite{W09}, rest frame NUV-optical SEDs \citep{W05} and $\mu$ m data from Spitzer \citep{G09}." +". In hat series of papers. ""dusty red Llate-types? ancl ""optically psive late-tvpes” are [ound to be largely the same thing. with a non-zero (but significantly lowered) star formation rate revealed by the Li data."," In that series of papers, “dusty red late-types"" and “optically passive late-types"" are found to be largely the same thing, with a non-zero (but significantly lowered) star formation rate revealed by the IR data." + Μος. spirals/late types have been studied: in. several recent papers (Leectal.2008:Dene2009:Hughes&Cortese2009:&Llughes2009)... as well as Mahajan&DHavehaudhuryv(2009) who talk about blue passive ealaxies galaxies with blue colours. but showing no indication of recent. star formation in their spectra) which mostly appear to have late type morphologies and have very recently shut down star formation.," Red spirals/late types have been studied in several recent papers \citep{Lee08,D09,HC09,CH09}, as well as \citet{MR09} who talk about blue passive galaxies galaxies with blue colours, but showing no indication of recent star formation in their spectra) which mostly appear to have late type morphologies and have very recently shut down star formation." + l'hese might be the ogenitors of the red. spirals., These might be the progenitors of the red spirals. + DBundsetal.(2010). have μαuclicd the redshift evolution of red sequence galaxies with lisk like components in COSMOS ancl use it to estimate iab as many as of spiral galaxies must pass through us phase on the way to the red sequence - making it an --mportant evolutionary step., \citet{Bundy09} have studied the redshift evolution of red sequence galaxies with disk like components in COSMOS and use it to estimate that as many as of spiral galaxies must pass through this phase on the way to the red sequence - making it an important evolutionary step. + The Galaxy Zoo project (Lintottetal.2008). revealed 10 presence of a significant number of visually classified spiral galaxies which are redder than the blue cloud (between of the tota galaxy population depending on environment. Bamlordetal. 2009)).," The Galaxy Zoo project \citep{L08} revealed the presence of a significant number of visually classified spiral galaxies which are redder than the blue cloud (between of the total galaxy population depending on environment, \citealt{B09}) )." + In this paper we study in more detail the physical properties ancl environments of this population of red spiral galaxies., In this paper we study in more detail the physical properties and environments of this population of red spiral galaxies. + Galaxies drawn from. this population have the morphological appearance of spiral galaxies with a distinct spiral armi structure. but have rest-[rame colours which are as red as à typical elliptical galaxy. indicating little or no recent star formation activity.," Galaxies drawn from this population have the morphological appearance of spiral galaxies with a distinct spiral arm structure, but have rest-frame colours which are as red as a typical elliptical galaxy, indicating little or no recent star formation activity." + We are studying these objects in order to identifv the physical process which is most important in their formation., We are studying these objects in order to identify the physical process which is most important in their formation. + lt ds clear that all spiral galaxies can be allectecl by various physical. processes as they evolve in this paper we attempt to identify which are most important. for red spirals. asking how they are able to shut down star formation while retaining their spiral morphology.," It is clear that all spiral galaxies can be affected by various physical processes as they evolve – in this paper we attempt to identify which are most important for red spirals, asking how they are able to shut down star formation while retaining their spiral morphology." + A list of possible mechanisms includes processes that depend on environment such as: (1) galaxy-ealaxy interactions: in high density regions there is an increased probability of interaction with other galaxies., A list of possible mechanisms includes processes that depend on environment such as: (1) galaxy-galaxy interactions: in high density regions there is an increased probability of interaction with other galaxies. + Most major mergers destroy spiral structure (Yoomre&Toomre1972). unless they involve very gas rich progenitors (Llopkinsetal.2000).. but some interactions can be quite gentle Walker. Mihos Llerneauist 1996). for example minor-mergers. tidal interactions etc. (," Most major mergers destroy spiral structure \citep{TT72} unless they involve very gas rich progenitors \citep{H09}, but some interactions can be quite gentle Walker, Mihos Hernquist 1996), for example minor-mergers, tidal interactions etc. (" +2) Interaction with the cluster itself also occurs and. can remove the gas which forms the reservoir for star formation.,2) Interaction with the cluster itself also occurs and can remove the gas which forms the reservoir for star formation. + This can be due to tidal ellects (c.g.Cinedin.2003)... or interaction with the hot intercluster gas. cither through thermal evaporation (Cowie&Songaila1977) or ram pressure stripping (Gunn&Gottτο). (," This can be due to tidal effects \citep[e.g.][]{Gn03}, or interaction with the hot intercluster gas, either through thermal evaporation \citep{CS77} or ram pressure stripping \citep{GG72}. (" +3). Processes like harrassment (Mooreetal.L999) ancl starvation or strangulation (Larsonetal.1980:Bekki2002) have also been shown to have a significant elfect on late type ealaxies.,"3) Processes like harrassment \citep{M99} and starvation or strangulation \citep{L80,Bekki02} have also been shown to have a significant effect on late type galaxies." + Llarrassment refers to the heating of gas by many small interactions. while starvation or strangulation refers to the eradual. exhaustion of disk eas after the hot halo has been stripped away.," Harrassment refers to the heating of gas by many small interactions, while starvation or strangulation refers to the gradual exhaustion of disk gas after the hot halo has been stripped away." + These mechanisms both occur at much larger. cluster radii lower densities) than the “elassice” environmental effects.," These mechanisms both occur at much larger cluster radii lower densities) than the ""classic"" environmental effects." + Internal. mechanisms could be more important., Internal mechanisms could be more important. + For example. (4) the latest," For example, (4) the latest" +were found.,were found. + The profiles of variability. ave dilferent. from. season to season., The profiles of variability are different from season to season. + The mean magnitude in V -band remained the same (liu=10.439 mag) during 20022008. while the amplitude: decreased. abruptly in. 2005.," The mean magnitude in $V$ -band remained the same $V_{\rm mean} = +10.439$ mag) during 2002–2008, while the amplitude decreased abruptly in 2005." + The. proposed interpretation is a redistribution of spots/spot groups over the surface of the star. while the total percentage of the spotted: area. was assumed to be constant within the error limits.," The proposed interpretation is a redistribution of spots/spot groups over the surface of the star, while the total percentage of the spotted area was assumed to be constant within the error limits." + A detailed. periodogram study of the photometric data enabled: us to derive a more accurate value for. the period of FR Cne., A detailed periodogram study of the photometric data enabled us to derive a more accurate value for the period of FR Cnc. +" We find. that P?0.08265+0.000015 d. In addition. we also presented DVRL, photometric calibration of 166 stars in FR Cne vicinity. whose V-maenituce is in the range of 0.85. 18.06 mag."," We find that $P = 0.08265 +\pm 0.000015$ d. In addition, we also presented $BVR_{c}I_{c}$ photometric calibration of 166 stars in FR Cnc vicinity, whose $V$ -magnitude is in the range of $9.85$ $18.06$ mag." + The DV? broad-band polarimetric observations of ER Cre have been obtained at ARLES in Nainital (Lucia) at Alanora Peak., The $BVR$ broad-band polarimetric observations of FR Cnc have been obtained at ARIES in Nainital (India) at Manora Peak. + Ehe observed polarization in D-band is well matched with the theoretical. values expected for Zeeman polarization model., The observed polarization in $B$ -band is well matched with the theoretical values expected for Zeeman polarization model. + However. the observed. polarization in V and 7? bands slightly exceeds the theoretical values and Thompson and Ravleigh scattering. from. inhomogeneous regions are not enough to explain the observed. polarization excess.," However, the observed polarization in $V$ and $R$ bands slightly exceeds the theoretical values and Thompson and Rayleigh scattering from inhomogeneous regions are not enough to explain the observed polarization excess." + Therefore the excess of linear. polarization should come from an additional source of polarization., Therefore the excess of linear polarization should come from an additional source of polarization. + Taking into account that we conclude that FR Che is not a binary. the mechanism which can produce additional linear polarization is probably scattering in circumsellar material distributed in an asymmetric geometry (o.&. see Pandevetal... 20093).," Taking into account that we conclude that FR Cnc is not a binary, the mechanism which can produce additional linear polarization is probably scattering in circumstellar material distributed in an asymmetric geometry (e.g. see \citeauthor{Pandey09}, \citeyear{Pandey09}) )." + A total of 58 spectra of FR Che. which have. been obtained in 20042008. were analvsed in this work.," A total of 58 spectra of FR Cnc, which have been obtained in 2004–2008, were analysed in this work." + Based on our spectroscopic observations. FR Cne was classified as," Based on our spectroscopic observations, FR Cnc was classified as" +withLST for the ACS Vireo Cluster Survey as well as or a GO program.,with for the ACS Virgo Cluster Survey as well as for a GO program. +" The trausieut is not associated with any star-fornune region aud the absolute magnitude of he progenitor is £ünter than M,=1 (67 MM... no correcting for extinction: Ofeketal.2008)).", The transient is not associated with any star-forming region and the absolute magnitude of the progenitor is fainter than $M_{g} \approx -4$ $<$ $_{\odot}$ not correcting for extinction; \citealt{okr+08}) ). + Thus. a nassive-star origin is quite uulikelv.," Thus, a massive-star origin is quite unlikely." + Iu contrast. 220088. 3300-OT..— iuc PTFl1lOfqs— occured| iu star-forming galaxies.," In contrast, 2008S, 300-OT, and 10fqs occurred in star-forming galaxies." + It may ο worth noting here that three supernovae (all of the core-collapse variety) have previously been discoverec in the host ealaxv of LiOfqs., It may be worth noting here that three supernovae (all of the core-collapse variety) have previously been discovered in the host galaxy of 10fqs. + It is perhaps of some sienificance that eight supernovae (six core-collapse. two unuclassified) were discovered iu 66916 in addition to 220088. Ouly oue supernova (of Type Τα) has Όσοι discovered in 3300.," It is perhaps of some significance that eight supernovae (six core-collapse, two unclassified) were discovered in 6946 in addition to 2008S. Only one supernova (of Type Ia) has been discovered in 300." + Sinall-uunber statistics and discovery bias (iucoupleteness from variety of different searches} notwithstanding. we ialke the sueecstion that galaxies with a high supernova rate pretercutially produce huninous red novae.," Small-number statistics and discovery bias (incompleteness from variety of different searches) notwithstanding, we make the suggestion that galaxies with a high supernova rate preferentially produce luminous red novae." + If this sugeestion is correct. then it would be worth the effort to svstematically maintain close vigilance on the nearest galaxies haviug laree supernova rates.," If this suggestion is correct, then it would be worth the effort to systematically maintain close vigilance on the nearest galaxies having large supernova rates." + Ίνκαetal.2007 suggested that MMon. V1332 Ser and M31 RV av— also be Iuninous. red novac.," \citealt{kor+07} suggested that Mon, V4332 Sgr and M31 RV may also be luminous, red novae." + We note here that the two Galactic sources are located in star-forming regions., We note here that the two Galactic sources are located in star-forming regions. +" Specifically, MMon is ina voung MMsyr) star cluster and may even have a B3 colmpanion (Afsar&Bond2007)."," Specifically, Mon is in a young Myr) star cluster and may even have a B3 companion \citep{ab07}." +. SSer (Alartinietal.1999) 18 located towards the iuuer Calasy (in Sagittarius)., Sgr \citep{mwt+99} is located towards the inner Galaxy (in Sagittarius). + On the other hand. M31 RV is located in the bulee of ALG3L.," On the other hand, M31 RV is located in the bulge of M31." +ZEST observations (undertakeu with WEPC?2 in parallel iode) taken about a decade ago show that the immediate envirous of M3I-RV are typical bulec-population stars (Boud&Siegel2006)., observations (undertaken with WFPC2 in parallel mode) taken about a decade ago show that the immediate environs of M31-RV are typical bulge-population stars \citep{bs06}. +. No usual reniunant star is seen at the astrometric position of M31 RV. nor any evidence of a Light echo (cousisteut with the absence of deuse circumstellar or interstellar eas that is essential to form echoes).," No unusual remnant star is seen at the astrometric position of M31 RV, nor any evidence of a light echo (consistent with the absence of dense circumstellar or interstellar gas that is essential to form echoes)." + Separately. there is uo evidence for any bhDunuiunous outbursts iu this area in the period 19121993 (Boschi&Minar2001).," Separately, there is no evidence for any luminous outbursts in this area in the period 1942–1993 \citep{bm04}." +. Thus. ΑΟ]. RV appears to have been a cataclysmic event in the bulee of M31.," Thus, M31 RV appears to have been a cataclysmic event in the bulge of M31." + PTFllOfqs is the fourth inenber of a class of extragalactic whicl possess a peak huninosity between that of novae and supernovae. aud have spectral aud photometric evolution that bear no resciublance to either supernovae or novae.," 10fqs is the fourth member of a class of extragalactic which possess a peak luminosity between that of novae and supernovae, and have spectral and photometric evolution that bear no resemblance to either supernovae or novae." + The other members of this class are MBS5-OT. 3300-OT. 220088. 3300-0T aid 220088 are remarkable for their very bright uud-iutrared proge1utors.," The other members of this class are M85-OT, 300-OT, 2008S. 300-OT and 2008S are remarkable for their very bright mid-infrared progenitors." + Though scusitive pre-explosiou observations of VS5-OT aud 110Éqs do exist. the large distance to the Vireo Cluster MMpce) relative to that of 3300) AIAIpec} and NGC69L6 MMpe) restIts in weak constraints ou the Iuninosityv of any. pre-explosion star.," Though sensitive pre-explosion observations of M85-OT and 10fqs do exist, the large distance to the Virgo Cluster Mpc) relative to that of 300 Mpc) and NGC6946 Mpc) results in weak constraints on the luminosity of any pre-explosion star." + 110fqs. 3300-OT. and 220088 occured in star-forming regions whereas ALS85-OT was in the bulee.f," 10fqs, 300-OT, and 2008S occurred in star-forming regions whereas M85-OT was in the bulge.," +acie. this eroup of explosive events cau be divided iuto a disk xd a bulge eroup., this group of explosive events can be divided into a disk and a bulge group. + The discovery of LiOfqs im itself cannot address whether the two eroups of huminous. red πονας are one iid the same.," The discovery of 10fqs in itself cannot address whether the two groups of luminous, red novae are one and the same." + The proposed models to explain this eroup are diverse: clectrou capture within an extreme asvinptotic giant brauch. (AGB) star. common-cuvelope phase (stellar mereer). inspiral of a elaut planet iuto the euvelope of au aging pareut star. a niost peculiar nova. and a most peculiar supernova.," The proposed models to explain this group are diverse: electron capture within an extreme asymptotic giant branch (AGB) star, common-envelope phase (stellar merger), inspiral of a giant planet into the envelope of an aging parent star, a most peculiar nova, and a most peculiar supernova." + The possible evidence of the broad feature centered around the Ca TW near-IR triplet with an interred velocity spread of kii t may be an important clue., The possible evidence of the broad feature centered around the Ca II near-IR triplet with an inferred velocity spread of km $^{-1}$ may be an important clue. + It would mean that these eveuts possess both a low- and a high-velocity outflow., It would mean that these events possess both a low- and a high-velocity outflow. + By comparison with other astronomical sources. one can envisage a high-velocity polar outflow aud a slower equatorial outflow (but with a larger mass).," By comparison with other astronomical sources, one can envisage a high-velocity polar outflow and a slower equatorial outflow (but with a larger mass)." + To this eud. continued seusitive spectroscopy of Ταν (and. of course other such future events) would be very valuable.," To this end, continued sensitive spectroscopy of 10fqs (and of course other such future events) would be very valuable." +" The ""Tranusieuts in the Local Universe” key project of the Palomar Transicut Factory is cesigued to systematically unveil events iu the gap between novae and supernovac.", The “Transients in the Local Universe” key project of the Palomar Transient Factory is designed to systematically unveil events in the gap between novae and supernovae. + It surveys 620.000 nearby galaxies (d<<200NIMpc) vearly at l-dav cadence auc a depth of Ro<21 mag. (," It surveys $\approx$ 20,000 nearby galaxies $d < 200$ Mpc) yearly at 1-day cadence and a depth of $R < 21$ mag. (" +"If the maxiuun luminosity of this class is) τας, then we would be seusitive to events out to LOOAINIpe.)","If the maximum luminosity of this class is $-14$ mag, then we would be sensitive to events out to Mpc.)" + Furthermore. Spifzer has a erowing archive of deep images of nearby galaxies (c.e.. SINGS. Iwenmicuttetal.2003: LVL. Daleetal.2009.. aud SiG. Shethotal. 2010)). andWYSE (Wirieltctal. has an ongoie all-sky survev in the mid-IR.," Furthermore, has a growing archive of deep images of nearby galaxies (e.g., SINGS, \citealt{kab+03}; LVL, \citealt{dcj+09}, and S4G, \citealt{srh+10}) ), and \citep{wem+10} has an ongoing all-sky survey in the mid-IR." + This will allow us to probe deeper iu search of the pre-explosion counterpart and possibly preseut a new chanucl for discovery of luminous red novae., This will allow us to probe deeper in search of the pre-explosion counterpart and possibly present a new channel for discovery of luminous red novae. + The discoverv of 1Ll0fqs is oulv the harbinger of the uucovering of a large sample of such trausieuts to unveil the nature of this new class of explosions., The discovery of 10fqs is only the harbinger of the uncovering of a large sample of such transients to unveil the nature of this new class of explosions. + ALALÉE. thanks the Cordon aud Betty Aloore Foundation for a Wale Fellowship/ im support of graduate study., M.M.K. thanks the Gordon and Betty Moore Foundation for a Hale Fellowship in support of graduate study. + The Weizimaun Institute PTF participation is supported in part by the Israel Science. Foundation via evauts to ACY., The Weizmann Institute PTF participation is supported in part by the Israel Science Foundation via grants to AGY. + The Weizmann-Caltech collaborative PTF effort i$ supported bv the US-Ixaecl Dinatioual Science Foundation., The Weizmann-Caltech collaborative PTF effort is supported by the US-Israel Binational Science Foundation. +" ACY and MS ave jointly supported by the ""making connections Weizuiuin-Uk program.", AGY and MS are jointly supported by the “making connections” Weizmann-UK program. +" ACY further acknowledges support bv a Marie Curie IRG fellowship and the Peter aud Patricia Cuuber Award. as well as funding bv the Benozivo Center for Astrophysics and the Yeda-Sela center at the Weizmann Tustitute,"," AGY further acknowledges support by a Marie Curie IRG fellowship and the Peter and Patricia Gruber Award, as well as funding by the Benoziyo Center for Astrophysics and the Yeda-Sela center at the Weizmann Institute." + A.V.F.s eroip and KAIT are supported by National Science Fouudution (NSF) eraut AST-0908886. the Svlvia Jim [ataman Foundation. the Richard Rhoda Coldiuan Fuid. Cary and Cyuthia Beueier. aud the TADASGO Foiudation: additional fuuding was provided bv NASA thaoughSpitzer erant 1322321. as well asHST exaut. AR-11218 from the Space Telescope Science Tustitute. which is operated by Associated Universities for Research in Απομών Tne... uuder NASA coutract NAS 5-26555.," A.V.F.'s group and KAIT are supported by National Science Foundation (NSF) grant AST-0908886, the Sylvia Jim Katzman Foundation, the Richard Rhoda Goldman Fund, Gary and Cynthia Bengier, and the TABASGO Foundation; additional funding was provided by NASA through grant 1322321, as well as grant AR-11248 from the Space Telescope Science Institute, which is operated by Associated Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555." + J.S.B. and his group are partially funded by a DOE SciDAC eraut., J.S.B. and his group are partially funded by a DOE SciDAC grant. + E.O.O. and D.P. are supported by the Einstein fellowship., E.O.O. and D.P. are supported by the Einstein fellowship. + L.D. issupported by the National Science. Foundation uuder, L.B. issupported by the National Science Foundation under +This upper limit of >=0.12 was selected: so as to niaximise the sample size vet minimise the error introduced » the isophotal corrections.,This upper limit of $z=0.12$ was selected so as to maximise the sample size yet minimise the error introduced by the isophotal corrections. + At 2=0.12 the uncertainty in. the isophotal. correction. (5:755).LU due to tvpe uncertaintyH (sce Appendix A). remains smaller than the photometric error (40.2 mags).," At $z=0.12$ the uncertainty in the isophotal correction $\pm^{0.07}_{0.16}$ ), due to type uncertainty (see Appendix A), remains smaller than the photometric error $\pm 0.2$ mags)." + Note that the increase in the error in the isophotal correction is primarily because of the increase in hein/rinsie isophotal limit due to a combination of surface xiehtness-dimming and the Ix-correction., Note that the increase in the error in the isophotal correction is primarily because of the increase in the isophotal limit due to a combination of surface brightness-dimming and the K-correction. + The final sample is therefore. pseudo: colieme-limited and contains 20.765 ealaxics. with recshilts. selected from a xwent sample of 45.000.," The final sample is therefore pseudo -limited and contains 20,765 galaxies, with redshifts, selected from a parent sample of 45,000." + Note: all magnitude and surface brightnesses are in the APM b; filter., Note: all magnitude and surface brightnesses are in the APM $b_{j}$ filter. + The APM magnitudes have already been corrected assuming a Ciaussian profile (see Maddox 1990b for Pull details)., The APM magnitudes have already been corrected assuming a Gaussian profile (see Maddox 1990b for full details). + This was aimed. primarily at recovering the light lost due to the seeing and is crucial for compact objects., This was aimed primarily at recovering the light lost due to the seeing and is crucial for compact objects. + Ht is known to significantIy underestimate the isophotal correction required [or low surface brightness disks., It is known to significantly underestimate the isophotal correction required for low surface brightness disks. + Such systems. typically. exhibit exponential profiles with disks which can extend a substantial distance bevond the isophote. the most. famous example being Malin 1 (Bothun 1987).," Such systems typically exhibit exponential profiles with disks which can extend a substantial distance beyond the isophote, the most famous example being Malin 1 (Bothun 1987)." + Once thought ol as a Virgo dwarf this svstem remains the most luminous Ποιά galaxy known., Once thought of as a Virgo dwarf this system remains the most luminous field galaxy known. + Yo compliment the Gaussian correction (required. for compact objects but inelfectual for extended: sources) we introduce an additional correction (inelfectual for compact sources but suitable for extended. clisks)., To compliment the Gaussian correction (required for compact objects but ineffectual for extended sources) we introduce an additional correction (ineffectual for compact sources but suitable for extended disks). + This. correction assumes all objects can be represented by a pure exponential surface brightness profile extending from the core outwards., This correction assumes all objects can be represented by a pure exponential surface brightness profile extending from the core outwards. +" In this case the surface brightness profile is simply: Or. Where X, is the central surface brightness in Wm arcsecτι a ds the scale length of the galaxy in aresees and r the racius in arcsecs."," In this case the surface brightness profile is simply: or, Where $\Sigma_o$ is the central surface brightness in $^{-2}$ $^{-2}$, $\alpha$ is the scale length of the galaxy in arcsecs and r the radius in arcsecs." + (A is the central surface brightness in. mag 2, $\mu_o$ is the central surface brightness in mag $^{-2}$. + Under this assumption a galaxys observed. isophotal luminosity is the integrated radial profileout to ris., Under this assumption a galaxy's observed isophotal luminosity is the integrated radial profileout to $r_{iso}$. +" which can be expressed in magnitudes as: (here qn"" denotes the apparent surface brightness uncorrected for redshift.) pia,"," which can be expressed in magnitudes as: (here $\mu^{app}_{o}$ denotes the apparent surface brightness uncorrected for redshift.) $\mu_{lim}$," +" the cleteetion/photometry isophote. can be expressed as: Aes Miso. Vise ancl qnn, are directly measurable quantities. equations (10 11 ) can be solved numerically."," the detection/photometry isophote, can be expressed as: As $m_{iso}$, $r_{iso}$ and $\mu_{lim}$ are directly measurable quantities, equations ( \ref{eq:miso} \ref{eq:mulim} ) can be solved numerically." + The total magnitude is then given by: Or. From this description an cxtrapolated central surface rightness can be deduced: numerically from. the specified isophotal area and isophotal magnitude (after the sccing correction).," The total magnitude is then given by: or, From this description an extrapolated central surface brightness can be deduced numerically from the specified isophotal area and isophotal magnitude (after the seeing correction)." + Note that this prescription ignores the possible oesence of a bulge. opacity. ancl inclination leading to an uneerestimate of the isophotal correction.," Note that this prescription ignores the possible presence of a bulge, opacity, and inclination leading to an underestimate of the isophotal correction." + This is unavoidable as the data quality is insullicient. to establish xilge-to-disk ratios., This is unavoidable as the data quality is insufficient to establish bulge-to-disk ratios. + Lo verify the impact of this we explore he accuracy of the isophotal correction for a variety. of galaxy types in Appendix A. The tests show that the isophotal correction is a significant improvement over the isophotal magnitudes for all types - apart for. cllipticals where the introduced. error is negligible compared. to the photometrie error - and a dramatic improvement for low surface brightness systems., To verify the impact of this we explore the accuracy of the isophotal correction for a variety of galaxy types in Appendix A. The tests show that the isophotal correction is a significant improvement over the isophotal magnitudes for all types - apart for ellipticals where the introduced error is negligible compared to the photometric error - and a dramatic improvement for low surface brightness systems. + The final magnitudes. after isophotal correction. now lie well within the quoted. error of £0.2 mags for both high- and Low-surface brightness ealaxies.," The final magnitudes, after isophotal correction, now lie well within the quoted error of $\pm 0.2$ mags for both high- and low-surface brightness galaxies." + Alost results cited in the literature use the central surface brightness or the cllective surface brightness., Most results cited in the literature use the central surface brightness or the effective surface brightness. + The central surface brightness. as described above. is the extrapolated surface brightness at the core under the assumption of a perfect. exponential disk.," The central surface brightness, as described above, is the extrapolated surface brightness at the core under the assumption of a perfect exponential disk." + The effective surface brightness is the mean surface brightness within the half-light radius., The effective surface brightness is the mean surface brightness within the half-light radius. + Phe conversion between the measures is relatively straightforward and described as follows: which can be solved numerically to ect The ellective surface brightness is now given by: Lenee from the isophotal magnitudes and areas we can derive the total magnitude anc effective surface brightness (quantities which are now independent of the isophotal detection threshold)., The conversion between the measures is relatively straightforward and described as follows: which can be solved numerically to get The effective surface brightness is now given by: Hence from the isophotal magnitudes and areas we can derive the total magnitude and effective surface brightness (quantities which are now independent of the isophotal detection threshold). + We chose to work with elfective surface brightness as it can. at some later stage. be measured directly. from higher quality CCD cata.," We chose to work with effective surface brightness as it can, at some later stage, be measured directly from higher quality CCD data." + Note that these surface brightness measures are allapperce/ rather thaninbrinsic. however this is not important as although surface brightness is distance dependent the isophotal correction is not (this is because both ιν and fn. vary. with redshift in the same wav).," Note that these surface brightness measures are all rather than, however this is not important as although surface brightness is distance dependent the isophotal correction is not (this is because both $\mu_{lim}$ and $\mu_{e}$ vary with redshift in the same way)." + Fie., Fig. + 2 shows the final 2dkFGRS sample after isophotal correction) for those galaxies with (upper panel), 2 shows the final 2dFGRS sample after isophotal correction) for those galaxies with (upper panel) +where Lush is the hybrid. svnchrotron. luminosity. and CDUtho qs the normalization. constant [for. thermal Comptonization spectrum of hybrid svnchrotron photons.,where $L_{\rm S}^{\rm nth}$ is the hybrid synchrotron luminosity and ${\cal C}^{\rm nth}$ is the normalization constant for thermal Comptonization spectrum of hybrid synchrotron photons. + ‘To caleulate the ratio we assume the optically thick hybrid svachrotron emission below a4 to be xSMOr). the source function of power-law electrons. where asPhe emissionM above £P] hn:ds xviplige and then we have The above expressions were derived. using the emission and absorption coellicients. for power-law electrons. lutivbicki Lightman 1979).," To calculate the ratio we assume the optically thick hybrid synchrotron emission below $\xnth$ to be $\propto S^{\rm pl}(x)$, the source function of power-law electrons, where The emission above $\nnth$ is $\propto \nu^{-(p-1)/2}$ and then we have The above expressions were derived using the emission and absorption coefficients for power-law electrons Rybicki Lightman 1979)." +" Since we neglected. the cut-olf in the electron distribution. the above relation holdsfor ⋎↓⊳↓⊳∢⊾⊳∖∖⋎∖∖⊽↓⊔⋅⊔↿↓↥⋖⋅≼∼∪⊔⇂↓⋅↓∣⋡⋯↓∪⊔⇂↓⋅∪⊔↓∢⊾↓⋖⋅≼∙↿↓⋅∪⊔≱∖il2o. . ⋅ 52. c5,to the emission attL aT with can be neglected."," Since we neglected the cut-off in the electron distribution, the above relation holdsfor $\xnth/x_{\rm c}\ll\gamma_{\rm f}^2$, when the contribution from electrons with $\gamma>\gamma_{\rm f}$ to the emission at $\xnth$ can be neglected." + The total luminosity of the CS process in the hybrid case is then When a«1. we can neglect LU and L2 in equations ancl respectively. since most of the luminosity is then produced at much higher energies. μιO.," The total luminosity of the CS process in the hybrid case is then When $\alpha<1$, we can neglect $L_{\rm S}^{\rm th}$ and $L_{\rm S}^{\rm nth}$ in equations and respectively, since most of the luminosity is then produced at much higher energies, $x \sim \Theta$." + Then. the ratio LPALYY can be caleulated as We can neglect the last term in the brackets if 6 is sullicientlv small. in which case Ll(C)«LP(cl).," Then, the ratio $L_{\rm CS}^{\rm nth}/L_{\rm CS}^{\rm th}$ can be calculated as We can neglect the last term in the brackets if $\delta$ is sufficiently small, in which case $L_{\rm C}^{\rm pl}\left({\cal C}^{\rm th}\right) \ll +L_{\rm C}^{\rm th}\left({\cal C}^{\rm th}\right)$." +" We then have Note that in caleulations of C""! *owe used the source function for power-Iaw. electrons.", We then have Note that in calculations of ${\cal C}^{\rm nth}$ we used the source function for power-law electrons. + Thus. this approximation does not have the correct thermal limit when 53iX1. Le. when tltzze (soe reffvidmabb. where a svnchrotron spectrum for such a case is shown).," Thus, this approximation does not have the correct thermal limit when $\gamma_{\rm nth} \gg \gamma_{\rm t}$, i.e. when $x_{\rm +t}^{\rm nth} \approx x_{\rm t}^{\rm th}$ (see \\ref{f:widma}b b, where a synchrotron spectrum for such a case is shown)." + Note also that for. harder power-law electron distributions (p& 3). for which the approximation does not hold. the amplification of CS luminosity will be even larger than that predicted above.," Note also that for harder power-law electron distributions $p\la3$ ), for which the approximation does not hold, the amplification of CS luminosity will be even larger than that predicted above." + We now calculate the ratio LP/LUS as a function of the energy. content of the electron power-law tail in three dilferent cases., We now calculate the ratio $L_{\rm CS}^{\rm nth}/L_{\rm CS}^{\rm th}$ as a function of the energy content of the electron power-law tail in three different cases. + They. correspond to the models of therma CS emission [rom a hot accretion dise with magnetic Lele in equipartition with the gas internal energy (dominated by hot ions) of WNZ00 for the Sevfert galaxy NGC 4151 ane two DIIDs in their hard states. GX 339.4 and (νο X-1.," They correspond to the models of thermal CS emission from a hot accretion disc with magnetic field in equipartition with the gas internal energy (dominated by hot ions) of WZ00 for the Seyfert galaxy NGC 4151 and two BHBs in their hard states, GX 339–4 and Cyg X-1." + The C'S luminosities predicted by the thermal mocdel SNa in WZOO) were orders of magnitude below the ones observe from NGC 4151 and GX 3394. and less bv a factor of ~3 for (νο N-1.," The CS luminosities predicted by the thermal model 8a in WZ00) were orders of magnitude below the ones observed from NGC 4151 and GX 339–4, and less by a factor of $\sim 3$ for Cyg X-1." +" Lor NGC 4151. we assume (see WZOO for details anc references) R=12.10 em. B2X8«107 Coa=OS5. Q—0.1. p24. «s,=100."," For NGC 4151, we assume (see WZ00 for details and references) $R=1.2 \times +10^{14}$ cm, $B=3.8 \times 10^3$ G, $\alpha=0.85$, $\Theta=0.1$, $p=4$, $\gamma_{\rm f}=100$." +" For GX 339.4. the same values of O. p and ; are adopted while 4?2SS10"" em. B= Goa2 0.75."," For GX 339–4, the same values of $\Theta$, $p$ and $\gamma_{\rm f}$ are adopted while $R=8.8 \times 10^{6}$ cm, $B=1.5 \times +10^7$ G, $\alpha=0.75$ ." +" In the case ofCvg N-1. P=3:10 em. B=10"" €. à=0.6. and we use the electron distribution obtained in McConnell et al. ("," In the case of Cyg X-1, $R=3 \times 10^{7}$ cm, $B=6.3 +\times 10^6$ G, $\alpha=0.6$, and we use the electron distribution obtained in McConnell et al. (" +20002) by fitting the spectra from the COMP'TIZL ancl OSSL detectors aboardCORO. O=017. tan=2.12. p=45. at assumed ο=107. which corresponds to a rather large value of à=0.19.,"2000a) by fitting the spectra from the COMPTEL and OSSE detectors aboard, $\Theta=0.17$, $\gamma_{\rm nth}=2.12$, $p=4.5$, at assumed $\gamma_{\rm f} =10^3$, which corresponds to a rather large value of $\delta=0.19$." + The results are shown in reff:powers.., The results are shown in \\ref{f:powers}. . + In the case of NGC 4151. and GNX 3394. à non-thermal tail in the electron distribution with 6<10.7 amplifies the power produced by the CS process sulliciently to reproduce the luminositiesof those sources.," In the case of NGC 4151 and GX 339–4, a non-thermal tail in the electron distribution with $\delta \la +10^{-3}$ amplifies the power produced by the CS process sufficiently to reproduce the luminositiesof those sources." + Therefore. unlike the results obtained. by WZO00 [rom investigating," Therefore, unlike the results obtained by WZ00 from investigating" +Josh Winn for helpful suggestions related to binary PAIS stars.,Josh Winn for helpful suggestions related to binary PMS stars. + We (hank Eric Feigelson Lor helpful discussions aud for his leadership of the COUP survey which produced the important ONC comparison sample of Fig., We thank Eric Feigelson for helpful discussions and for his leadership of the COUP survey which produced the important ONC comparison sample of Fig. + 2., 2. + This material is partly based on work supported by the National Aeronautics and Space Administration under Grant NACG5-12502 issued through the Origins of Solar Systems Program., This material is partly based on work supported by the National Aeronautics and Space Administration under Grant NAG5-12502 issued through the Origins of Solar Systems Program. +In the past two cecaces. astrophysical timing measurements has been used to infer the existence of multiple low-mass planetary objects.,"In the past two decades, astrophysical timing measurements has been used to infer the existence of multiple low-mass planetary objects." + Wolszezan&Frail(1992). announced the first detection of a planetary system around the pulsar PSR1257|12., \cite{WolszczanFrail1992} announced the first detection of a planetary system around the pulsar PSR1257+12. + A two-planet system orbiting the short-period subclwarl B of the eclipsing binary HIN Virginis was first. presented in the works bv Leeetal.(2009).. and recently. Deuermannetal. (2010)... Potteretal.(2011)... ane Dovleetal...(2011). announced. the existence of two cireumbinary planets in possible mean motion resonances around NN Serpentis. two giant planets orbiting the eclipsing binary UZ Fornacis. and a single cireumbinary. planet around the stars of the binary system Ixepler 16. respectively.," A two-planet system orbiting the short-period subdwarf B of the eclipsing binary HW Virginis was first presented in the works by \cite{Lee2009}, and recently \cite{Beuermann2010}, , \cite{Potter2011}, and \cite{Doyle2011} announced the existence of two circumbinary planets in possible mean motion resonances around NN Serpentis, two giant planets orbiting the eclipsing binary UZ Fornacis, and a single circumbinary planet around the stars of the binary system Kepler 16, respectively." + The studies by Beuermannetal.(2010). and Potteretal.(2011). inferred the presence of additional massive objects bv explaining the observed timing anomalies with the light-travel time (LET. hereafter). effect.," The studies by \cite{Beuermann2010} and \cite{Potter2011} + inferred the presence of additional massive objects by explaining the observed timing anomalies with the light-travel time (LTT, hereafter) effect." + In. the ideal case. the stellar components of an cclipsing binary system are orbiting their common centre of mass with a constant period.," In the ideal case, the stellar components of an eclipsing binary system are orbiting their common centre of mass with a constant period." + Timing irregularities can occur if the binary svstem is accompanied by an additional massive object., Timing irregularities can occur if the binary system is accompanied by an additional massive object. + In this case. the binary centre of mass is orbiting the svstem's centre of mass.," In this case, the binary centre of mass is orbiting the system's centre of mass." + At some times the binary will be closer to the observer while at other times. it will be farther away. giving rise to the LIT elfect on measured eclipse egress times.," At some times the binary will be closer to the observer while at other times, it will be farther away, giving rise to the LTT effect on measured eclipse egress times." + In à recent studs. the measurements of eclipse egress times of the eclipsing polar binary HEU Aquarii (MU «qr. hereafter) was used to infer the presence of a cireumbinary. planet around this svstem (Qianetal.2011).," In a recent study, the measurements of eclipse egress times of the eclipsing polar binary HU Aquarii (HU Aqr, hereafter) was used to infer the presence of a circumbinary planet around this system \citep{Qian2011}." +. These authors modeled the complete timing data set by adding the LPT clleets from two cireumbinary. planets and found that the pericentre of the outer planetary companion is inside the orbit of the inner planet., These authors modeled the complete timing data set by adding the LTT effects from two circumbinary planets and found that the pericentre of the outer planetary companion is inside the orbit of the inner planet. + Phis orbital architecture implies a crossing orbit configuration which points to strong mutual perturbations between the two planets., This orbital architecture implies a crossing orbit configuration which points to strong mutual perturbations between the two planets. + Horneretal.(2011) subsequently carried out a detailed stability analvsis of these bodies and found that almost all their initial conditions lead to unstable orbits on short-time scales., \cite{Horner2011} subsequently carried out a detailed stability analysis of these bodies and found that almost all their initial conditions lead to unstable orbits on short-time scales. + These authors concluded that the HU Aqr planetary svstem has most likely a clillerent orbital architecture than proposed by Qianetal.(2011)., These authors concluded that the HU Aqr planetary system has most likely a different orbital architecture than proposed by \cite{Qian2011}. + Three possibilities exist to explain the instability of the proposed. planetary system., Three possibilities exist to explain the instability of the proposed planetary system. + Either i) the LIT. parameter space was not explored thoroughly while orbital stability constrains were imposed. omitted inQianetal. (2011)]]. or. ii) the applied," Either i) the LTT parameter space was not explored thoroughly while orbital stability constrains were imposed [omitted in\cite{Qian2011}] ], or, ii) the applied" +While it is recognized that long gamma-ray bursts (GRBs) are associated with very energetic supernovae. sometime referred to as hypernovae. our understanding of the physics that drives and connects these events is far from being established.,"While it is recognized that long gamma-ray bursts (GRBs) are associated with very energetic supernovae, sometime referred to as hypernovae, our understanding of the physics that drives and connects these events is far from being established." + Observationally. long GRBs are more complicated than were thought before.," Observationally, long GRBs are more complicated than were thought before." + For example. they appear to have subgroup that includes GRB 980425 etal. 1998).. GRB 031203 (Malesant.etal.2004;Soderberg and GRB 060218 (Campanaetal.2006;Cobbal.2006;Pianet 2006).," For example, they appear to have subgroup that includes GRB 980425 \citep{Galama1998}, GRB 031203 \citep{Malesani2004,Soderberg2004} and GRB 060218 \citep{Campana2006,Cobb2006,Pian2006}." +. These GRBs occurred relatively nearby and their energies radiated by prompt gamma rays were significantly smaller than the other long GRBs. and they were associated with well-studied hypernovae (SN 1998bw. SN 2003Iw and SN 2006a). respectively).," These GRBs occurred relatively nearby and their energies radiated by prompt gamma rays were significantly smaller than the other long GRBs, and they were associated with well-studied hypernovae (SN 1998bw, SN 2003lw and SN 2006aj, respectively)." + Radio observations of these events (Kulkarnietal.1998:Soderberg2006) suggest the presence of mildly relativistic ejecta. which ts a different component from the usual nonrelativistic component of the supernova explosion.," Radio observations of these events \citep{Kulkarni1998,Soderberg2006} suggest the presence of mildly relativistic ejecta, which is a different component from the usual nonrelativistic component of the supernova explosion." + Their rate of occurrence may be an order of magnitude higher than that of the more energetic GRBs (Liangetal.2007:Soderberg2006).," Their rate of occurrence may be an order of magnitude higher than that of the more energetic GRBs \citep{Liang2007,Soderberg2006}." +. See also. e.g.. Liangetal.(2006):Murase(2006);TomaZhang(2007) for other followup studies.," See also, e.g., \citet*{Liang2006,Murase2006,Toma2007,Waxman2007,Gupta2007} for other followup studies." + Very recently. à mildly relativistic outflow component has also been inferred in à supernova of type Ibe unassociated with a GRB. SN 2008D (Soderbergetal.2008).," Very recently, a mildly relativistic outflow component has also been inferred in a supernova of type Ibc unassociated with a GRB, SN 2008D \citep{Soderberg2008}." +. In this paper. we investigate the broadband radiation from the mildly relativistic ejecta associated with hypernovae. focusing especially on the high-energy photon emission in the X-ray and gamma-ray ranges.," In this paper, we investigate the broadband radiation from the mildly relativistic ejecta associated with hypernovae, focusing especially on the high-energy photon emission in the X-ray and gamma-ray ranges." + We mainly study the radiation from relativistic electrons which are primarily accelerated in shocks. but we also consider the radiation from secondary electrons which are generated by interactions involving accelerated protons.," We mainly study the radiation from relativistic electrons which are primarily accelerated in shocks, but we also consider the radiation from secondary electrons which are generated by interactions involving accelerated protons." + The latter were extensively studied by Asano&Mészaros(2008) in connection with the origin of Galactic cosmic rays (Wangetal.2007;Budnik 2008).," The latter were extensively studied by \citet{Asano2008} in connection with the origin of Galactic cosmic rays \citep{Wang2007,Budnik2008}." +. Our treatment of the proton component includes also the previously neglected effect of photopair production., Our treatment of the proton component includes also the previously neglected effect of photopair production. + We show that the primary electron component photon signature generally dominates over the proton component. for a wide range of energies.," We show that the primary electron component photon signature generally dominates over the proton component, for a wide range of energies." + In particular the inverse-Compton scattering of hypernova photons due to the primary electrons appears the most promising channel for detection in the X-ray and gamma-ray ranges., In particular the inverse-Compton scattering of hypernova photons due to the primary electrons appears the most promising channel for detection in the X-ray and gamma-ray ranges. + The paper is organized as follows., The paper is organized as follows. + In 2.. we introduce the relevant supernova parameters and the spectrum of the primary accelerated electrons.," In \ref{sec:Spectrum of Primary Electrons}, we introduce the relevant supernova parameters and the spectrum of the primary accelerated electrons." + We present our main results on the radiation from the primary electrons in 3.. and discuss its dependence on the parameters.," We present our main results on the radiation from the primary electrons in \ref{sec:Radiation from Primary Electrons}, and discuss its dependence on the parameters." + Section + is devoted to a discussion of the photon radiation of à proton origin and its comparison to the electron component.," Section \ref{sec:Radiation from secondary electrons and proton +acceleration} is devoted to a discussion of the photon radiation of a proton origin and its comparison to the electron component." + We summarize our conclusions in 5.., We summarize our conclusions in \ref{sec:Conclusions}. . +or foreground signal) nicely reflects the underlying distribution of magnetic field within the large scale structure.,or foreground signal) nicely reflects the underlying distribution of magnetic field within the large scale structure. + However. a very large number of lines of sight probed by RM measurements (much more than the 3072 used in this investigation) are needed to overcome the statistical noise induced by the particular magnetic field realization within the cosmic structures. in order to distinguish between the wide range of models we used here.," However, a very large number of lines of sight probed by RM measurements (much more than the 3072 used in this investigation) are needed to overcome the statistical noise induced by the particular magnetic field realization within the cosmic structures, in order to distinguish between the wide range of models we used here." + In general. the RM signal is strongly dominated by the denser regions (e.g. those populated by galaxy clusters and groups) and not by the low density ones. like filaments.," In general, the RM signal is strongly dominated by the denser regions (e.g. those populated by galaxy clusters and groups) and not by the low density ones, like filaments." + On this point. the magnetic field associated with filaments already changes by several orders of magnitudes within the different models used here.," On this point, the magnetic field associated with filaments already changes by several orders of magnitudes within the different models used here." + Aditionally. the normalized estimator iyi; 1s extremely sensitive to measurement errors and to the presence of the GF (despite attempts to remove it by subtracting a smoothed map).," Aditionally, the normalized estimator $\omega_{RM}$ is extremely sensitive to measurement errors and to the presence of the GF (despite attempts to remove it by subtracting a smoothed map)." + It is fair to say that given the current measurement errors in the available RMs and our knowledge of the GF. present studies cannot determine the magnetization magnitude of the Universe based only on the cross-correlation wyay. Whatever the significance of the measured signal is.," It is fair to say that given the current measurement errors in the available RMs and our knowledge of the GF, present studies cannot determine the magnetization magnitude of the Universe based only on the cross-correlation $\omega_{RM}$, whatever the significance of the measured signal is." + On the contrary. the shape of the unnormalized estimator £j; (the same as used by Lee et al.," On the contrary, the shape of the unnormalized estimator $\xi_{RM}$ (the same as used by Lee et al." + 2009) is relatively insensitive against the presence of measurement errors for the RMs and for the presence of the GF (as long as the described removal technique is used)., 2009) is relatively insensitive against the presence of measurement errors for the RMs and for the presence of the GF (as long as the described removal technique is used). + Its amplitude. however. is quite strongly affected by measurement uncertainties.," Its amplitude, however, is quite strongly affected by measurement uncertainties." + Current measurement errors (as for example those inherited by the Taylor's published sample) suppress the signal by a significant amount in such a way that it is impossible to relate the amplitude of the cross-correlation function to the underlying magnetization of the the large scale structure., Current measurement errors (as for example those inherited by the Taylor's published sample) suppress the signal by a significant amount in such a way that it is impossible to relate the amplitude of the cross-correlation function to the underlying magnetization of the the large scale structure. + However. we expect that future radio telescopes will be able of reaching error magnitudes of order of | rad 7 that could make the correction of the signal possible.," However, we expect that future radio telescopes will be able of reaching error magnitudes of order of 1 rad $^{-2}$ that could make the correction of the signal possible." + Unfortunately. this estimator does not nicely encode in its shape the details of the magnetization of the large scale structure and. especially. its amplitude is extremely sensitive to the of the Universe.," Unfortunately, this estimator does not nicely encode in its shape the details of the magnetization of the large scale structure and, especially, its amplitude is extremely sensitive to the of the Universe." + Therefore. any interpretation of an observed signal is limited by our knowledge of the redshift distribution of the sources (towards the RM signals measured). as well as by our knowledge of the distribution and evolution of the cosmie universal magnetization.," Therefore, any interpretation of an observed signal is limited by our knowledge of the redshift distribution of the sources (towards the RM signals measured), as well as by our knowledge of the distribution and evolution of the cosmic universal magnetization." + Future observational data will help to put better constraints on theoretical models for the origin of cosmological magnetic fields which. in return. can be implemented in next generation of MHD cosmological simulations in order to draw a self-consistent picture that can be compared against observations.," Future observational data will help to put better constraints on theoretical models for the origin of cosmological magnetic fields which, in return, can be implemented in next generation of MHD cosmological simulations in order to draw a self-consistent picture that can be compared against observations." +>=1.221. and is very similar to ours. with strong OL] emission. and. some absorption features. inclucing (weak) Balmer lines. (,"$z=1.221$, and is very similar to ours, with strong [OII] emission and some absorption features, including (weak) Balmer lines. (" +Qi) Object 1412 is GDS J033235.78-214621.5.,iii) Object 1412 is GDS J033235.78-274627.5. + The FORS2 spectrum gave z=1.094. consistent with our z=1.092. and ike ours shows OL] emission.," The FORS2 spectrum gave $z=1.094$, consistent with our $z=1.092$, and like ours shows [OII] emission." + The LIS and L9 lines we find are not identified on the FORS2 spectrum. but it extends urther into the red to reveal 7/72 and £20 absorption.," The H8 and H9 lines we find are not identified on the FORS2 spectrum, but it extends further into the red to reveal $H\gamma$ and $H\delta$ absorption." + Εις on the basis of both spectra this seems to be a strong Balmer absorption galaxy. (, Thus on the basis of both spectra this seems to be a strong Balmer absorption galaxy. ( +(iv) Object 1777 was not observed in the FORS2 survey. (,iv) Object 1777 was not observed in the FORS2 survey. ( +(v) Object 2536 is GDS J033231.83-214356.2.,v) Object 2536 is GDS J033231.83-274356.2. + Phe FORS2 spectrum gave z=1.550. slightly higher than our στ= 1.513.," The FORS2 spectrum gave $z=1.550$, slightly higher than our $z=1.513$ ." + Like our GALOS spectrum it features Megll (2800) sorption. and extends a little further recwards. detecting an Oll] emission linc. (," Like our GMOS spectrum it features MgII $2800\rm \AA$ ) absorption, and extends a little further redwards, detecting an [OII] emission line. (" +(vi) Object 1064 is CDS J033230.02-274726.8.,vi) Object 1064 is GDS J033230.02-274726.8. + Phe FORS2 spectrum gave z= 1.553. a close match to our z=154s. and detects. Mell (2800) absorption. and a weak OL) emission line at the reddest extreme. ((vii£viifix) Objects 1143 ancl 1565 and 2401 were not observed in the FORS2 survey. (," The FORS2 spectrum gave $z=1.553$ , a close match to our $z=1.548$, and detects MgII $2800\rm \AA$ ) absorption, and a weak [OII] emission line at the reddest extreme, (vii/viii/ix) Objects 1143 and 1565 and 2401 were not observed in the FORS2 survey. (" +(x) Object 1876 is CDS J033230.34-274523.6.,x) Object 1876 is GDS J033230.34-274523.6. + Phe FORS2 spectrum gave z= 1.223. slightlv higher than our 2= 1.196.," The FORS2 spectrum gave $z=1.223$ , slightly higher than our $z=1.196$ ." + The FORS2 spectrum. like ours. shows Mgll (2800A) absorption. OL] emission. and. Balmer lines. extending further rechwares to ffs.," The FORS2 spectrum, like ours, shows MgII $2800\rm \AA$ ) absorption, [OII] emission and Balmer lines, extending further redwards to $H\gamma$ ." + 1876 is also number 2158 in the Doherty ct al. (, 1876 is also number 2158 in the Doherty et al. ( +2005) DEIAIOS sample. where it has 2=1.222 and Uh. 16 absorption. very weak OLI]. anc is classed as I|e ((xi) Object 2246 is CDS J033229.82-274510.8.,"2005) DEIMOS sample, where it has $z=1.222$ and $\rm H\&K$, $\rm H\delta$ absorption, very weak [OII], and is classed as $\rm E+e$ (xi) Object 2246 is GDS J033229.82-274510.8." + Phe FORS2 spectrum gave z=1.225. slightly greater than our 2= 1.196.," The FORS2 spectrum gave $z=1.225$, slightly greater than our $z=1.196$ ." + Like our spectrum it shows OL] emission and. H&Ix absorption. the spectra are very similar but ours has better signal-to-noise. (," Like our spectrum it shows [OII] emission and $\rm H\&K$ absorption, the spectra are very similar but ours has better signal-to-noise. (" +(xü/xii) Objects 1823 and 1955 were not observed in the FORS2 survey. (,xii/xiii) Objects 1823 and 1955 were not observed in the FORS2 survey. ( +(xiv) Object 1762 is CDS J033227.02-274407.2.,xiv) Object 1762 is GDS J033227.02-274407.2. + The FORS2 spectrum gave z=1.127. slightlv higher than our z=1.102. and like ours shows moderate OL] emission and. weak Balmer absorption.," The FORS2 spectrum gave $z=1.127$, slightly higher than our $z=1.102$, and like ours shows moderate [OII] emission and weak Balmer absorption." + The authors mark £75 but it is barely visible. (, The authors mark $H\gamma$ but it is barely visible. ( +"xv) Objects 1843. and 1158 were not observed in the 0m""OBS2 survey.",xv) Objects 1843 and 1158 were not observed in the FORS2 survey. +" However. the AGN 1843 was observed with a""LTE FORS in an earlier survey of Chandra sources."," However, the AGN 1843 was observed with VLT FORS in an earlier survey of Chandra sources." + Szokoly et al. (, Szokoly et al. ( +2004) initially proposed 2=(0.484 on the basis of a single unidentified line.,2004) initially proposed $z=0.484$ on the basis of a single unidentified line. + Zheng et al. (, Zheng et al. ( +2004) claimed >=1.59 on the basis of a single line. this time supported by photometric redshifts from. COMBDO-17. and. hyper-z.,"2004) claimed $z=1.89$ on the basis of a single line, this time supported by photometric redshifts from COMBO-17 and $z$." + Llowever. for our. 9-band. photometry we obtain a > solution of 2=2.06. and our continuum is significantly better fitted for z=2.017 than +=L89.," However, for our 9-band photometry we obtain a $z$ solution of $z=2.06$, and our continuum is significantly better fitted for $z=2.017$ than $z=1.89$." + Considering that photometric redshifts are uncertain by at least SA. we retain >=2.017 as à provisional and approximate estimate.," Considering that photometric redshifts are uncertain by at least $8\%$, we retain $z=2.017$ as a provisional and approximate estimate." + Near infra-red spectroscopy. should detect multiple emission lines and eive an unambiguous answer., Near infra-red spectroscopy should detect multiple emission lines and give an unambiguous answer. + In summary. out of our redshift sample of 16. SN objects were also observed by. Vanzella et al. (," In summary, out of our redshift sample of 16, 8 objects were also observed by Vanzella et al. (" +2005) and/or Doherty et al. (,2005) and/or Doherty et al. ( +2005). who obtained. very similar spectra. similar spectra (although in two cases with a redshift cliscrepancy that might be larger than expected from statistical error).,"2005), who obtained very similar spectra similar spectra (although in two cases with a redshift discrepancy that might be larger than expected from statistical error)." + Seven of our redshifts appear to be new., Seven of our redshifts appear to be new. + Finally. for the 202 AGN. more data is needed.," Finally, for the $z\sim 2$ AGN, more data is needed." + We fit our ERG spectra. at. 5800A«9000A where the signal-to-noise is greatest. with a series of 40 template spectra from Jimenez et al. (," We fit our ERG spectra, at $5800<\lambda<9000\rm \AA$ where the signal-to-noise is greatest, with a series of 40 template spectra from Jimenez et al. (" +2004).,2004). + Phese represent passively evolving stellar populations. with a Salpeter IME and solar metallicity. at ages from 0.001 to 14 Car.," These represent passively evolving stellar populations, with a Salpeter IMF and solar metallicity, at ages from 0.001 to 14 Gyr." + We experimented with models of lower anc higher metallicity. but found no indication that these fitted: any better., We experimented with models of lower and higher metallicity but found no indication that these fitted any better. + Furthermore the colours (Caputi ct al., Furthermore the colours (Caputi et al. + 2004) of our ERGs ancl the ultraviolet spectroscopy of de Mello et al. (, 2004) of our ERGs and the ultraviolet spectroscopy of de Mello et al. ( +2004) support the assumption that ERGs have approximately solar metallicity. even {ο το2.,"2004) support the assumption that ERGs have approximately solar metallicity, even to $z\sim 2$." + From the scatter between the 17 ‘blank’ (i.e. no signal) slit spectra. we derived a function representing the noise per pixel produced by the sky. background: within the area of a single spectroscopic slit (Fig 2b).," From the scatter between the 17 `blank' (i.e. no signal) slit spectra, we derived a function representing the noise per pixel produced by the sky background within the area of a single spectroscopic slit (Fig 2b)." + This is much worse at the red end due to strong sky emission lines and the declining detector response., This is much worse at the red end due to strong sky emission lines and the declining detector response. + We fit models to a combination of the GCMOS spectra and the 9-band (DVLZLEN. 3.6/4.5m) photometry. using a least-squares (minimum 47) method. with cachpixel of the spectrum. weighted: using this noise function.," We fit models to a combination of the GMOS spectra and the 9-band $BVIZJHK$, $3.6/4.5\mu \rm m$ ) photometry, using a least-squares (minimum $\chi^2$ ) method, with eachpixel of the spectrum weighted using this noise function." + The GALOS spectrum and the 9-band. photometry make similar contributions to to the combined 27. which is minimized in the fitting.," The GMOS spectrum and the 9-band photometry make similar contributions to to the combined $\chi^2$, which is minimized in the fitting." + An initial series of fits as a function of redshift provided the continuum-fit redshift σος for each ERG., An initial series of fits as a function of redshift provided the continuum-fit redshift $z_{cf}$ for each ERG. +" We then fixed the redshift at.this. value and. fit with aAeo-component model. this being the sum of a single-age passively evolving model. of age Z,,:. and a 'starburst component. of much vounger stars. represented by the Jimenez et al. ("," We then fixed the redshift atthis value and fit with a model, this being the sum of a single-age passively evolving model, of age $T_{pas}$, and a `starburst' component of much younger stars, represented by the Jimenez et al. (" +2004) model. integrated with a constant SER for a time interval 7.5 priorto the time of observation.,"2004) model, integrated with a constant SFR for a time interval $T_{sb}$ priorto the time of observation." + Lt is important to note that this component is not a single-age population. but contains stars of all ages from zero to 754.," It is important to note that this component is not a single-age population, but contains stars of all ages from zero to $T_{sb}$." + Phe voung component is subject to dust-reddening. following the extinction curve of Calzetti et al. (," The young component is subject to dust-reddening, following the extinction curve of Calzetti et al. (" +2000). parzmeterized in terms of (5.V).,"2000), parameterized in terms of $E(B-V)$." +" The normalization of the voung component is parameterized as fu, the fraction of the emitted Dux it contributes in the rest-[rame D-band. 4500A."," The normalization of the young component is parameterized as $f_{sb}$ the fraction of the emitted flux it contributes in the rest-frame $B$ -band, $4500\rm \AA$ ." + Llenee it is described by three parameters. ανν E(DV) and. fa.," Hence it is described by three parameters, $T_{sb}$, $E(B-V)$ and $f_{sb}$." +" For cach ERG. we searched for the best-fittingmoclel. by calculatinga 47for the GMOS spectrum. plus the 9-band photometry compared with a set of model spectra. forming a grid of 18000 points in the 4-dimensional (Z5,45. fs. 1.5. ἀπ V)) parameter space."," For each ERG, we searched for the best-fittingmodel, by calculatinga $\chi^2$for the GMOS spectrum, plus the 9-band photometry compared with a set of model spectra, forming a grid of 18000 points in the 4-dimensional $T_{pas}$ , $f_{sb}$ , $T_{sb}$ , $E(B-V)$ ) parameter space." +" Z5, was allowed to vary [rom 10 toSOO Myr and E(D13 from Oto0.5 mag (ely~2 ", $T_{sb}$ was allowed to vary from 10 to800 Myr and $E(B-V)$ from 0 to0.5 mag $A_V\sim 2$ +that of the youngest haloes. bu the oldest haOOS are sienificanthy more clustered than the oldest haloes.,"that of the youngest haloes, but the oldest haloes are significantly more clustered than the oldest haloes." + The cllects become very large. for he lowest. masses tial we resolve., The effects become very large for the lowest masses that we resolve. + A lothTAL. the larec-scale. autocorrelajon amplitude for the oldest haloes is more than 5 times that for the voungest haloes., At $10^{11}h^{-1}{\rm M_\odot}$ the large-scale autocorrelation amplitude for the oldest haloes is more than 5 times that for the youngest haloes. + In Fig. 4..," In Fig. \ref{fig:visual}," +" we provide some images to give a visual impression of he relative distributions of voung"" and ""old"" haloes.", we provide some images to give a visual impression of the relative distributions of “young” and “old” haloes. + Here we show haloes with particle number in he range 100.200] in a slice through the Millennium Simulation 30h.+Alpe thick.," Here we show haloes with particle number in the range $[100, 200]$ in a slice through the Millennium Simulation $30 +h^{-1} {\rm Mpc}$ thick." + The top row shows the positions of the voungest of these haloes (Left). of the oklest (middle). and. of an equal number of dark matter particles selected at random within the slice (right).," The top row shows the positions of the youngest of these haloes (left), of the oldest (middle), and of an equal number of dark matter particles selected at random within the slice (right)." + Phe bottom row shows corresponding plots or the tails., The bottom row shows corresponding plots for the tails. + di is striking hat although the haloes (hy definition) avoid massive clumps in the dark matter distribution. the old haloes ollow the large-scale cosmic web quite closely. while the cisribution of voung haloes looks almost uniform.," It is striking that although the haloes (by definition) avoid massive clumps in the dark matter distribution, the old haloes follow the large-scale cosmic web quite closely, while the distribution of young haloes looks almost uniform." + lig., Fig. + 5 explores the formation time dependence of the clustering bias in more cleail., \ref{fig:biasinbin} explores the formation time dependence of the clustering bias in more detail. + We take the sample of all haloes with particle number in the range 100.200] and we split it into ten equal-sized subsamples by formation time.," We take the sample of all haloes with particle number in the range $[100, 200]$ and we split it into ten equal-sized subsamples by formation time." + We then compute bias factors and mean formation redshifts, We then compute bias factors and mean formation redshifts +properties.,properties. + Moreover. in the MNB and LNB spectra. the TC parameter are also weIl constrained. differently from what reported in DSO2 where they strongly deviate from a smooth trend with unreasonable associated errors.," Moreover, in the MNB and LNB spectra, the TC parameter are also well constrained, differently from what reported in DS02 where they strongly deviate from a smooth trend with unreasonable associated errors." + We claim that one reason of the DSO2 results could be the rough timing selection for spectral analysis (see Section 2)., We claim that one reason of the DS02 results could be the rough timing selection for spectral analysis (see Section 2). + The quantity that undergoes the major change as a function of the source position in the CD/HID is in fact the estimated 0.1—200 keV luminosity which has a maximum increase of about from the UHB to the UNB-MNB state., The quantity that undergoes the major change as a function of the source position in the CD/HID is in fact the estimated 0.1–200 keV luminosity which has a maximum increase of about from the UHB to the UNB-MNB state. + Less significant variations are observed in the plasma temperature AT). ard the thermal component (CAF). defined as the ratio of erergy fluxes of Comptonized over seed photons.," Less significant variations are observed in the plasma temperature $\kte$ and the thermal component (CAF), defined as the ratio of energy fluxes of Comptonized over seed photons." + The CAF. in fact. decreases monotonically fron the UHB to the LNB.," The CAF, in fact, decreases monotonically from the UHB to the LNB." + We also observe that the inferred plasma temperature. optical depth and luminosity in the LNB slightly deviates from the monotonic trend seen in the first four spectra.," We also observe that the inferred plasma temperature, optical depth and luminosity in the LNB slightly deviates from the monotonic trend seen in the first four spectra." + Whether this small jump ts real or due to some bias in the spectral modeling is not straightforward., Whether this small jump is real or due to some bias in the spectral modeling is not straightforward. + In particular the presence of a strong | keV and weaker 2.5 keV emission feature especially in the NB spectra may play some role in the actual determination of the temperature of the seed photons of the pure thermal mmodel. which is reflected in the AZ. and à values.," In particular the presence of a strong 1 keV and weaker 2.5 keV emission feature especially in the NB spectra may play some role in the actual determination of the temperature of the seed photons of the pure thermal model, which is reflected in the $\kte$ and $\alpha$ values." + The new theoretical and observational support provided in the last years (e.g.. TMK97. P06. F08) to the study of hard X-ray emission in LMXBs significantly pushed forward cur knowledge about the nature of the emission processes in these sources.," The new theoretical and observational support provided in the last years (e.g., TMK97, P06, F08) to the study of hard X-ray emission in LMXBs significantly pushed forward our knowledge about the nature of the emission processes in these sources." + Multi-wavelength observations of Z sources (e.g..22) strengthened the evidence that the motion of the sources along the Z track is mainly controlled by the accretion rate 1 which increases from the HB to the FB.," Multi-wavelength observations of Z sources \citep[e.g.,][]{penninx88, vrtilek90} strengthened the evidence that the motion of the sources along the Z track is mainly controlled by the accretion rate $\dot{M}$ which increases from the HB to the FB." + However. little is said on why an increase in A/ leads to a disappearance of the hard tail in LMXBs and. moreover. the actual meaning of AS itself is not fully understood.," However, little is said on why an increase in $\dot{M}$ leads to a disappearance of the hard tail in LMXBs and, moreover, the actual meaning of $\dot{M}$ itself is not fully understood." + It is however widely accepted that the simple source bolometric X-ray luminosity is not a good A-tracer. a hint occurring not only in NS but also in black hole (BH) systems (e.g..2)..," It is however widely accepted that the simple source bolometric X-ray luminosity is not a good $\dot{M}$ -tracer, a hint occurring not only in NS but also in black hole (BH) systems \citep[e.g.,][]{vdk01}." + We propose to actually define Af-values for a single source. one related to the aceretion disk Αι and another one (A44) to the so-called transition layer (TL). the region where the disk angular velocity deviates from its Keplerian rotation in order to adjust to the angular velocity of the slowly spinning NS.," We propose to actually define $\dot{M}$ -values for a single source, one related to the accretion disk $\dot{M}_{\rm disk}$ and another one $\dot{M}_{\rm tl}$ ) to the so-called transition layer (TL), the region where the disk angular velocity deviates from its Keplerian rotation in order to adjust to the angular velocity of the slowly spinning NS." + The reason for this splitting i5 that at very high accretion rate levels (close to the Eddington limit). the radiation pressure from the accretion disk may eject a significantfraction of the accreted matter. producing a powerful wind surrounding the system (?)..," The reason for this splitting is that at very high accretion rate levels (close to the Eddington limit), the radiation pressure from the accretion disk may eject a significantfraction of the accreted matter, producing a powerful wind surrounding the system \citep{bradshaw07}." +" In this case the mass flow coming from the innermost part of the disk through the TL to theNS surface is less than that arriving at the disk outer part. namely Afm Mane with M,xMaa ("," In this case the mass flow coming from the innermost part of the disk through the TL to theNS surface is less than that arriving at the disk outer part, namely $\dot{M}_{\rm tl} \la \dot{M}_{\rm disk}$ , with $\dot{M}_{\rm tl} \propto \dot{M}_{\rm disk}$ \citep{tsa07}. ." +2. , +is the most important. factor that determines the fragmentation Gimescale. (he spacing of cores in the filament. and the masses of cores (Fiege&Pudritz2000b).,"is the most important factor that determines the fragmentation timescale, the spacing of cores in the filament, and the masses of cores \citep{FP2}." +. Therefore. polarization observations of GI1.11-0.12. in combination with this work. would constrain the fragmentation properties of the cloud strongly. and determine the initial magnetohydrodyvnamic conditions that lead to star formation.," Therefore, polarization observations of G11.11-0.12, in combination with this work, would constrain the fragmentation properties of the cloud strongly, and determine the initial magnetohydrodynamic conditions that lead to star formation." + Figure 21 is the corresponding ligure for region 2., Figure \ref{fig:FP_F2_pol} is the corresponding figure for region 2. + All three tvpes of polarization patterns are possible for region 2., All three types of polarization patterns are possible for region 2. + The polarization percentages have a similar range of allowed values as for Region 1. although the scatter is much greater.," The polarization percentages have a similar range of allowed values as for Region 1, although the scatter is much greater." + This paper is the first. astronomical research problem attempted using a new genetic algorithm “Ferret” designed by one of the authors (J.F.)., This paper is the first astronomical research problem attempted using a new multi-objective genetic algorithm “Ferret” designed by one of the authors (J.F.). + We have tested (his code thoroughly. and it performs extremely well on data-fitiing problems where the goal is to [find theclass of allowed solutions. rather than a single representative solution.," We have tested this code thoroughly, and it performs extremely well on data-fitting problems where the goal is to find the of allowed solutions, rather than a single representative solution." + This allows the user to understand (he detailed structure of the parameter space. ancl gives an honest representation of the degeneracies of the models being tested.," This allows the user to understand the detailed structure of the parameter space, and gives an honest representation of the degeneracies of the models being tested." + Our previous analvsis of (he GI1.11-0.12 inlrared-dark cloud (Paper I) demonstrated that its density. profile is steeper than the r7 profile inferred [rom observations of other filamentary molecular elouds (Alves et al., Our previous analysis of the G11.11-0.12 infrared-dark cloud (Paper I) demonstrated that its density profile is steeper than the $~r^{-2}$ profile inferred from observations of other filamentary molecular clouds (Alves et al. + 1998. Lada οἱ al 1999. Johnstone Bally 1999). and that the profile can be fit reasonably well by the non-magnetic Ostriker model.," 1998, Lada et al 1999, Johnstone Bally 1999), and that the profile can be fit reasonably well by the non-magnetic Ostriker model." + This paper provides a more detailed analvsis (hat compares (he radial structure of the infrared-dark cloud with three models of filamentary clouds: the non-magnetic Ostriker model. the magnetic GS model. and the magnetic FP model.," This paper provides a more detailed analysis that compares the radial structure of the G11.11-0.12 infrared-dark cloud with three models of filamentary clouds: the non-magnetic Ostriker model, the magnetic GS model, and the magnetic FP model." + We found that the density profile of GI1.11-0.12 is consistent with the Ostriker model. but also with regimes of both magnetic models.," We found that the density profile of G11.11-0.12 is consistent with the Ostriker model, but also with regimes of both magnetic models." + Our technique allowed us to map out the allowed regions of parameter space (horoughiy for each model., Our technique allowed us to map out the allowed regions of parameter space thoroughly for each model. + Our most significant. result is that nearly all magnetic solutions that are in agreement wilh the cata are dominated by the poloidal component of the magnetic fiekl. or are magnetically neutral.," Our most significant result is that nearly all magnetic solutions that are in agreement with the data are dominated by the poloidal component of the magnetic field, or are magnetically neutral." +" A lew toroidal-dfield dominated solutions are allowed. but only for Region 2. which has a much “noisier” radial structure (han Region 1. and even (hese solutions are close to magnetic neutrality,"," A few toroidal-field dominated solutions are allowed, but only for Region 2, which has a much “noisier” radial structure than Region 1, and even these solutions are close to magnetic neutrality." + GI1.11-0.12 appears to be very dillerent [rom the filaments studied by FPL. which were inferred to have a shallow 7? density profile and a dominant toroidal field component.," G11.11-0.12 appears to be very different from the filaments studied by FP1, which were inferred to have a shallow $~r^{-2}$ density profile and a dominant toroidal field component." + Note that the FP model is the only model of the three models considered (hat can account for both toroidal-field dominated filaments and the, Note that the FP model is the only model of the three models considered that can account for both toroidal-field dominated filaments and the +this (wo-climensional model reduces (o a one-dimensional model which is presented in subsection ??..,this two-dimensional model reduces to a one-dimensional model which is presented in Subsection \ref{EM:mod_1D}. + Finally in Subsection ??.. we explain how to recover equation (1.1)) as a model for dislocation dvnamies with for some particular non-negative and svmmetric matrix 1.," Finally in Subsection \ref{EM:mod_1D_np}, we explain how to recover equation \ref{EM:burger}) ) as a model for dislocation dynamics with for some particular non-negative and symmetric matrix $A$." + We now present in details the 6wo-dimensional model., We now present in details the two-dimensional model. + We denote by A the vector X=(ry.-r2)€E2.," We denote by ${ + X}$ the vector ${ X} = (x_1,x_2) \in \R^2$." +" We consider a crystal filling the whole space E? and its displacement emοι,0):E?—E27. where we have not vet introduced the time We introduce the total strain z(0)=(z;(0));j-t2 which is à symmetric matrix delined by The total strain can be spittec in two parts: where 75;aC is the elastic strain and zTHAP is the plastic strain."," We consider a crystal filling the whole space $\R^2$ and its displacement $v=(v_1,v_2):\R^2\rightarrow\R^2$, where we have not yet introduced the time We introduce the total strain $\varepsilon(v)=(\varepsilon_{ij}(v))_{i,j=1,2}$ which is a symmetric matrix defined by The total strain can be spitted in two parts: where $\e^e_{ij}$ is the elastic strain and $\e^p_{ij}$ is the plastic strain." + The scalar function u is the −− ↕↽≻↥≀↧↪∖⋅⋃≺∢≺⇂↥⋟∖⋅↥↽≻↥≀↧↴≺∢≼↲∐∐↲↥∐≀↕⊔∖⋅⋟∖⋅∪≺∢↥≀↧↴∩↲≼⇂↥∪⋃∐↲∣⋅⊲−��↥⋟∖⋅∐↕↽≻⋟∖⋅∙∖⊐∖⋅∥↲∐↓∖∖⊽∐∪⋟∖⋅≼↲∐↓≀↧↴⋃⋅↥⇀↖⊆∶∣⋅∕↕⋝∖⋅≺⇂≼↲∐∐≼↲≺⇂∣↽≻∙∖⇁− − −∣∣⋡∕⋅− ∙ where (5.n)) is a family of vectors in I2. such that n js a unit vector orthogonal to the Burger's vector (see Hirth οἱ al.," The scalar function $u^k$ is the plastic displacement associated to the $k$ -th slip system whose matrix $\varepsilon^{0,k}_{ij}$ is defined by where $(b^{k}, n^{k})$ is a family of vectors in $\R^2$ , such that ${n}^{k}$ is a unit vector orthogonal to the Burger's vector ${b}^k$ (see Hirth et al." + |?|. for the definition of the Burger's vector of a To simplify (he presentation. we assume the simplest possible periodicitv property of the," \cite{HL92} for the definition of the Burger's vector of a To simplify the presentation, we assume the simplest possible periodicity property of the" +the directions for various scales seem to be raucomly scattered. aud the significance is very low.,"the directions for various scales seem to be randomly scattered, and the significance is very low." + Iu fact. only at 3° there seems to bea 26 hint. but auy exact directions are non-existeut.," In fact, only at $3^\circ$ there seems to be a $2 \sigma$ hint, but any exact directions are non-existent." + This is in clisagreciment with the results from Hausenetal.(2001).. who claimed that there exists a maxi of non-Caussianity on Lemispheres centered at the ecliptic poles.," This is in disagreement with the results from \cite{hansen:2004b}, who claimed that there exists a maximum of non-Gaussianity on hemispheres centered at the ecliptic poles." + However. when we ignore correlations and use a strictly diagonal covariance matrix. the results are more iu agreement with Hansenctal.(200L)..," However, when we ignore correlations and use a strictly diagonal covariance matrix, the results are more in agreement with \cite{hansen:2004b}." + These results are preseuted in Table 2.. where we note several deviations at the less than level.," These results are presented in Table \ref{tab:resultsdiagonal}, where we note several deviations at the less than $1\%$ level." + Also. when performing the directional analysis. we Πα that he nou-Cassian signal has a clear maxima close to (but not directly on) the ecliptic north-pole. as seen in fietre 6 and Table 2..," Also, when performing the directional analysis, we find that the non-Gaussian signal has a clear maximum close to (but not directly on) the ecliptic north-pole, as seen in figure \ref{fig:directions_old} and Table \ref{tab:resultsdiagonal}." + This is in agreement with the direction described by Hausenetal. (2001).., This is in agreement with the direction described by \cite{hansen:2004b}. +. As noted above. here we use a larger extended mask than in ITausenuetal.(200)..," As noted above, here we use a larger extended mask than in \cite{hansen:2004b}." + We have also tested with the smaller extended mash aud fud results siniar to the ones presented in Table 2.., We have also tested with the smaller extended mask and find results simiar to the ones presented in Table \ref{tab:resultsdiagonal}. + We conclude. that by using a diagonal approximaion to the correlation matrix. we still find non-Craussianities aud asviunetries," We conclude that by using a diagonal approximation to the correlation matrix, we still find non-Gaussianities and asymmetries" +the sky. polarisation. observations were mace alternately centred. on the photo Centre of the comet and on a region of the skv more than 30 arcmin away from the comet (along the anti-tail Polarisation standard 9 Geni was observed to calibrate the observed. position angle.,"the sky polarisation, observations were made alternately centred on the photo Centre of the comet and on a region of the sky more than 30 arcmin away from the comet (along the anti-tail Polarisation standard 9 Gem was observed to calibrate the observed position angle." + Comet’s ILI magnitudes were obtained using the observations of solar tvpe stars. namely 11D29461. LLD76151.," Comet's IHW magnitudes were obtained using the observations of solar type stars, namely HD29461, HD76151." + Polarisation values. corrected. position anele and ΗΝ magnitudes in continuum bands are given in Table 1..," Polarisation values, corrected position angle and IHW magnitudes in continuum bands are given in Table \ref{obstab}." + The various quantities obtained. from our observations on he comet τὸΠοιος are listed in Table 1. namely. the degree of (PX) and its error(c p). the position angle(8). brightness magnitudes. alone with observing ime(CJD'E). filter. aperture size. total integration time. and the orbital parameters at the time of observation.," The various quantities obtained from our observations on the comet 17P/Holmes are listed in Table \ref{obstab} namely, the degree of $P\%$ ) and its $\epsilon_{P}\%$ ), the position $\theta)$, brightness magnitudes along with observing time(JDT), filter, aperture size, total integration time, and the orbital parameters at the time of observation." + I he values Listed in Table 1. have been carefully. checked or any inconsistencyv., All the values listed in Table \ref{obstab} have been carefully checked for any inconsistency. + Figure 1. shows the LHleliocentric ancl Ceocentric ranges ancl the phase. angle(a) at. the ime of observation. which were obtained using the JPL’s," Figure \ref{17P_ephe} shows the Heliocentric and Geocentric ranges and the phase $\alpha$ ) at the time of observation, which were obtained using the JPL's" +oroperties of the average galaxy population. and with observational estimates.,"properties of the average galaxy population, and with observational estimates." + Fig., Fig. + + shows the K-band rest-frame luminosity function of ost galaxies for the HOST? ¢dot-dashed line) and the HOST3 (dashed line) samples. compared with the galaxy luminosity unction measured considering all galaxies in the simulation box. at different redshift.," \ref{fig:lumfunc} shows the K-band rest-frame luminosity function of host galaxies for the HOST2 (dot-dashed line) and the HOST3 (dashed line) samples, compared with the galaxy luminosity function measured considering all galaxies in the simulation box, at different redshift." + At all redshifts. LGRB host galaxies aave luminosities well below the characteristic luminosity L.. in agreement with observational measurements.," At all redshifts, LGRB host galaxies have luminosities well below the characteristic luminosity $L_*$, in agreement with observational measurements." + While the total uminosity function evolves strongly with the redshift (particularly beyond z~ 1). the number densities and the range of luminosities of LGRB host galaxies vary more mildly. due to the fact that at uigher redshift a larger fraction of the w10le galaxy population can vost LGRBs.," While the total luminosity function evolves strongly with the redshift (particularly beyond $z\sim 1$ ), the number densities and the range of luminosities of LGRB host galaxies vary more mildly, due to the fact that at higher redshift a larger fraction of the whole galaxy population can host LGRBs." + Recent observational studies have focused on the stellar mass distribution of GRB host galaxies., Recent observational studies have focused on the stellar mass distribution of GRB host galaxies. + ?. have found that the typical stellar mass of host galaxies is smaller than the stellar mass of field galaxies at the same redshift., \citet{CastroCeron_etal_2008} have found that the typical stellar mass of host galaxies is smaller than the stellar mass of field galaxies at the same redshift. +" For a sample of 30 LGRB hosts. they provide estimates of the stellar mass between LO"" and 1013AJ.. with a mean value of AZ,~10A7..."," For a sample of $30$ LGRB hosts, they provide estimates of the stellar mass between $10^7$ and $10^{11} \, M_{\odot}$, with a mean value of $M_*\sim10^{9.7}\,M_{\odot}$." + About 70 per cent of the host galaxies in their sample have stellar mass A.10472ALL.," About $70$ per cent of the host galaxies in their sample have stellar mass $M_*<10^{10.1}\, M_{\odot}$." + Similar results have been found by ?.., Similar results have been found by \citet{Savaglio_etal_2008}. + Using a sample of 46 GRB hosts -the largest sample so far- they estimatePa a median stelar muss of 10775AZ. and find that about 83 per cent of the studied systems have stellar mass between LO“? and 101775AL...," Using a sample of 46 GRB hosts -the largest sample so far- they estimate a median stellar mass of $10^{9.3}\,M_{\odot}$, and find that about $83$ per cent of the studied systems have stellar mass between $10^{8.5}$ and $10^{10.3}\,M_{\odot}$." + In Fig., In Fig. + 5. we compare the galaxy mass distribution for model host galaxies from the HOST? and HOSTS samples to the distribution obtained considering all galaxies in the simulated box., \ref{fig:massfunc} we compare the galaxy mass distribution for model host galaxies from the HOST2 and HOST3 samples to the distribution obtained considering all galaxies in the simulated box. + For this figure. all galaxies and hosts at all redshifts up to 2—9 have been used.," For this figure, all galaxies and hosts at all redshifts up to $z\sim 9$ have been used." + Fig., Fig. + 5. shows that LGRB host galaxies have typically low mass with a small fraction of them having stelar mass up to 10/5AZ.," \ref{fig:massfunc} shows that LGRB host galaxies have typically low mass with a small fraction of them having stellar mass up to $\sim +10^{11}\,M_{\odot}$." +" About 90 per cent of the host galaxies have stellar mass «10""AZ. and «107Al. for HOST3 and HOST? respectively."," About 90 per cent of the host galaxies have stellar mass $<10^{9}\,M_{\odot}$ and $<10^{10}\,M_{\odot}$ for HOST3 and HOST2 respectively." + It is well known that the galaxy stellar mass is tightly correlated with the rest-frame K-band luminosity., It is well known that the galaxy stellar mass is tightly correlated with the rest-frame K-band luminosity. +" ?. have shown that this relation applies to GRB host galaxies as well. but they argue that GRB galaxies have on average higher luminosity than ""feld"" galaxies with the same stellar mass. implying a lower Al,/Liy ratio, as expected for younger galaxies."," \citet{Savaglio_etal_2008} have shown that this relation applies to GRB host galaxies as well, but they argue that GRB galaxies have on average higher luminosity than “field” galaxies with the same stellar mass, implying a lower $M_*/L_{\rm K}$ ratio, as expected for younger galaxies." + We compare results from our model to observational measurements in Fig. 6.., We compare results from our model to observational measurements in Fig. \ref{fig:mass_colour}. + The dashed black line is the best-tit to the observational data by ?:: logAl;=0.463Aly 0.102.," The dashed black line is the best-fit to the observational data by \citet{Savaglio_etal_2008}: ${\rm log}\,M_* = +-0.463\times M_K-0.102$ ." + The red line in Fig., The red line in Fig. + 6 shows the mean luminosity-mass. relation— obtained by using all galaxies in the simulation boxes up to ς~9., \ref{fig:mass_colour} shows the mean luminosity-mass relation obtained by using all galaxies in the simulation boxes up to $z\sim 9$. + The pink line shows the mean value for host galaxies in the HOST2 sample. and the blue line corresponds to the mean value obtained for the HOST3 sample.," The pink line shows the mean value for host galaxies in the HOST2 sample, and the blue line corresponds to the mean value obtained for the HOST3 sample." + To compute the average of Adj: we weigh each host by the likelihood that it contains a GRB., To compute the average of $M_K$ we weigh each host by the likelihood that it contains a GRB. + For this sample. we also show the quartiles of the distribution.," For this sample, we also show the quartiles of the distribution." + We note that ? adopt a Baldry Glazebrook IMF for their stellar mass estimates. while the model used in this study adopts a Chabrier IMF to compute model magnitudes.," We note that \citet{Savaglio_etal_2008} adopt a Baldry Glazebrook IMF for their stellar mass estimates, while the model used in this study adopts a Chabrier IMF to compute model magnitudes." + In order to compare model results with observational estimates. we have decreased the observed stellar mass by a factor 1.3.," In order to compare model results with observational estimates, we have decreased the observed stellar mass by a factor $1.3$." + Fig., Fig. + 6 shows that the K-band absolute magnitude distribution of simulated GRB host galaxies is in good agreement with observations., \ref{fig:mass_colour} shows that the K-band absolute magnitude distribution of simulated GRB host galaxies is in good agreement with observations. +" It also shows that. on average. host galaxies have stellar masses which are lower. although with a large scatter. than ""typical"" galaxies wihthe same mass. in agreement with observational findings."," It also shows that, on average, host galaxies have stellar masses which are lower, although with a large scatter, than “typical” galaxies with the same mass, in agreement with observational findings." + In Fig. 7..," In Fig. \ref{fig:colour}," + we compare the median colour model LGRB host galaxies with observational measurements by of?.shownasblack symbols..," we compare the median colour of model LGRB host galaxies with observational measurements by \citet[][shown as black +symbols]{Savaglio_etal_2008}." + Model results indicate that GRB galaxies are typically bluer than the average galaxy population at the same redshift., Model results indicate that GRB galaxies are typically bluer than the average galaxy population at the same redshift. + We note that the observed colours exhibit a quite large scatter. probably due to the unknown dust extinction.," We note that the observed colours exhibit a quite large scatter, probably due to the unknown dust extinction." + Fig., Fig. + 8 shows the median gas metallicity evolution for the HOSTS sample (black crosses) compared with the observational estimates for the GRB-DLAs studied in Savaglio et al. (, \ref{fig:metalz} shows the median gas metallicity evolution for the HOST3 sample (black crosses) compared with the observational estimates for the GRB-DLAs studied in Savaglio et al. ( +2006. 2008 - blue triangles with error bars a few show the lower and upper-branch metallicity solution in Savaglio e al.,"2006, 2008 - blue triangles with error bars a few show the lower and upper-branch metallicity solution in Savaglio et al." + 2008)., 2008). + In order to compare with observations we weigh each host with the total number of LGRBs., In order to compare with observations we weigh each host with the total number of LGRBs. + The red line in Fig., The red line in Fig. + 8 shows the median metallicity obtained using all galaxies in the simulation box., \ref{fig:metalz} shows the median metallicity obtained using all galaxies in the simulation box. +" The figure shows that the metallicity of model galaxies (both ""normal"" and hosts) does not evolve significantly with redshift.", The figure shows that the metallicity of model galaxies (both “normal” and hosts) does not evolve significantly with redshift. + The observational measurements exhibit a large scatter and have typically large uncertainties., The observational measurements exhibit a large scatter and have typically large uncertainties. + Within these. model predictions are in relatively good agreement with observational data.," Within these, model predictions are in relatively good agreement with observational data." + It should, It should +"Equation 1.. Thus. for example. the fraction of «Me dwarls at V—£e = 1.0 is directly proportional to the relative utunber of stars which have formed in the last ~310* years. while fyay, at V—£e: = 2.6 correspouds to the relative number of stars with ages of less than 1 Gyr.","Equation \ref{equation-age-vi}, Thus, for example, the fraction of dMe dwarfs at $V-I_C$ = 1.0 is directly proportional to the relative number of stars which have formed in the last $\sim 3 \times 10^7$ years, while $f_{dMe}$ at $V-I_C$ = 2.6 corresponds to the relative number of stars with ages of less than 1 Gyr." + Iu this manuer. we cau determine fare lor a range of V—fe: colour. aud map the cumulative star formation history of the Galactic disk.," In this manner, we can determine $f_{dMe}$ for a range of $V-I_C$ colour, and map the cumulative star formation history of the Galactic disk." + We have applied tle analysis techuique outlined in the previous section to data for the M dwarfs in the VC sample., We have applied the analysis technique outlined in the previous section to data for the M dwarfs in the VC sample. + For stars lacking V.—Fe: ineasuremeuts. we use the TiO» index to estimate the colour.," For stars lacking $V-I_C$ measurements, we use the TiO5 index to estimate the colour." + The results are plotted in Figure 12.. giving [μις as a function of the inferred age.," The results are plotted in Figure \ref{fig-sfhist}, giving $f_{dMe}$ as a function of the inferred age." + A constant star formation rate (Le. fuga; directly proportional to age) is shown for reference as the thick solid line with slope unity in this figure., A constant star formation rate (i.e. $f_{dMe}$ directly proportional to age) is shown for reference as the thick solid line with slope unity in this figure. + Two clifferent ways of computing the fraction of dMe stars are illustrated., Two different ways of computing the fraction of dMe stars are illustrated. + The long-daslied line (connecting solid triaugles) gives the Traction found by welgine all stars equally., The long-dashed line (connecting solid triangles) gives the fraction found by weighting all stars equally. + The thin solid line (connecting open triaugles) slows the fraction found yw weighting each star bv the inverse of its W velocity. as proposed by Wielen(1971.1977).," The thin solid line (connecting open triangles) shows the fraction found by weighting each star by the inverse of its W velocity, as proposed by \citet{w74,w77}." +. The atte “approach allows fo “the increased scale height of -higher velocity (older) stars. aud Cousequent shorter resideuce time iu the Solar Neighbourhood.," The latter approach allows for the increased scale height of higher velocity (older) stars, and consequent shorter residence time in the Solar Neighbourhood." + However. this methocl adds statistical nolse Npacing slenificaut (tidue?)," However, this method adds statistical noise by placing significant (undue?)" + weight ou a simall uumber of Ligh velocity stars., weight on a small number of high velocity stars. + Iu both cases we exclde the SBI aud SB2 (short period binary) systems. since the activity tn those stars may be influenced by other effecs than age.," In both cases we exclude the SB1 and SB2 (short period binary) systems, since the activity in those stars may be influenced by other effects than age." + Figure 12. indicates hat the overall star formation history is broadly consistent with a constaut star formation rate., Figure \ref{fig-sfhist} indicates that the overall star formation history is broadly consistent with a constant star formation rate. + The major feature notable in the distribution is a step at ~ 1 Gyr which. taken at [ace value. would iudicate that 105€ of the local clisk stars formed in a burst at that time.," The major feature notable in the distribution is a step at $\sim$ 1 Gyr which, if taken at face value, would indicate that $\sim10$ of the local disk stars formed in a burst at that time." + However. this feature is unlikely to be real. since it corresponds with the Group B/D trausition i Figure 2. which we discussed at length in Section [.1..," However, this feature is unlikely to be real, since it corresponds with the Group B/D transition in Figure \ref{fig-tio5ha1} which we discussed at length in Section \ref{activity}." + Such a distiuct change in activity propertieW. is unlisely to be well described by a simple linear relation., Such a distinct change in activity properties is unlikely to be well described by a simple linear relation. + Data for clusters with ages betwee the Hyades (0.6 Cyr) and M67 (1 Cyr) would be useful in confirming whether this feature las a astroplivsical. rather than evolutionary. origin.," Data for clusters with ages between the Hyades (0.6 Gyr) and M67 (4 Gyr) would be useful in confirming whether this feature has an astrophysical, rather than evolutionary, origin." + Tve other characteristic of the star formation history illustrated in Figure 12. is a slight cleLicjency in the number oL young («I Gyr) stars relative to tlie number of older (21 Cyr) stars., The other characteristic of the star formation history illustrated in Figure \ref{fig-sfhist} is a slight deficiency in the number of young $<1$ Gyr) stars relative to the number of older $>1$ Gyr) stars. + At agesDm of less than 1 Gyr. both the weightede and unweighted[we solutions match the expectation of coustant star formation (slope unity). with slopes of 0.02£0.16 aud 0.95+0.50 respectively.," At ages of less than 1 Gyr, both the weighted and unweighted solutions match the expectation of constant star formation (slope unity), with slopes of $0.92 \pm 0.16$ and $0.95 \pm 0.50$ respectively." + For the full :vee distribution. however. we derive best-fit slopes of 1.1140.09 for the uuweighted solution aud 1.38£0.19 [or the weighted solution.," For the full age distribution, however, we derive best-fit slopes of $1.14 \pm 0.09$ for the unweighted solution and $1.38 \pm 0.19$ for the weighted solution." + This result implies a slight decrease in the star formation yate in recent times., This result implies a slight decrease in the star formation rate in recent times. + However. this could reflect iucompleteuess iu the VC sample which. siuce it is," However, this could reflect incompleteness in the VC sample which, since it is" + Comes(2003) 30735 (Morbidelhetal.2008).., \citet{Gomes03} $-$ \citep{SSBNMorbiAJ}. + been proposed that could move even cold KBO objects from a disk truncated at ~ 35 AU to their present distances (Levisonctal.2008)., been proposed that could move even cold KBO objects from a disk truncated at $\sim$ 35 AU to their present distances \citep{LevisonNice08}. +. Iu this model. the dvuanically cold KBO objects would have formed in the outer regions of tlhe massive protoplanetary disk. perhaps in the region around 30 AU.," In this model, the dynamically cold KBO objects would have formed in the outer regions of the massive protoplanetary disk, perhaps in the region around 30 AU." + The vast majoritv of the known cvuamically cold classical belt objects have semimajor axes )) between 12 and ls AU., The vast majority of the known dynamically cold classical belt objects have semimajor axes ) between 42 and 48 AU. + There are also low e low classical objects with « 39.1 AU.," There are also low , low classical objects with $<$ 39.4 AU." + This senmumajor axis is less than the current position of he Neptune 3:2 resonance. the semünajor axis at which Pluto aud the plutino IKBOs are currently ound.," This semimajor axis is less than the current position of the Neptune 3:2 resonance, the semimajor axis at which Pluto and the plutino KBOs are currently found." + The population of these objects. known as inner classical belt KBOs (ICIDBOs). with )etween about 36 and 39.1 AU. is much lower hau that of the main classical belt objects ( jetweecn 12 and Is AU).," The population of these objects, known as inner classical belt KBOs (ICKBOs), with between about 36 and 39.4 AU, is much lower than that of the main classical belt objects ( between 42 and 48 AU)." + Because of the relative dearth of known classical belt objects interior to he Neptune 3:2 resonance. no published plysical studies have specifically targeted these objects.," Because of the relative dearth of known classical belt objects interior to the Neptune 3:2 resonance, no published physical studies have specifically targeted these objects." + This class of objects extends the «range of objects compared to the rauge spanned by the main belt alone., This class of objects extends the range of objects compared to the range spanned by the main belt alone. + Thus it is potentially useful iu the study of auv eradicuts of chemical or surface processing history with formation distance for manor bodies in the outer solar system., Thus it is potentially useful in the study of any gradients of chemical or surface processing history with formation distance for minor bodies in the outer solar system. + Tere we present new optical colors of a sample of immer classical IKBOs., Here we present new optical colors of a sample of inner classical KBOs. + We compare the color distribution of these objects with that of other classesof outersolar svstcmobjects., We compare the color distribution of these objects with that of other classesof outersolar systemobjects. +"rather modest. of the order of 1—3xLO""M./year (Longetal.2003a). on cay 7 into oulburst (Table 1).","rather modest, of the order of $1-3 \times 10^{-9}M_{\odot}/year$ \citep{lon03a} on day 7 into outburst (Table 1)." + During the dip. the white dwarf might have a temperature of about 25. 000Ix (Ixniggeοἱal.2002) on dav 27.," During the dip, the white dwarf might have a temperature of about $25,000$ K \citep{kni02} on day 27." + This is more (han 10.000A. above ils quiescence temperature which is 2214.500A.," This is more than $10,000K$ above its quiescence temperature which is $\approx 14,500K$." + During the rebrightening phase. the mass accretion rate seenis (o be somewhat smaller than during the plateau and only peaks at LxI09M. ! (Longetal.2003a) on clay 46.," During the rebrightening phase, the mass accretion rate seems to be somewhat smaller than during the plateau and only peaks at $1\times 10^{-9} M_{\odot}$ $^{-1}$ \citep{lon03a} on day 46." + We put these values into the table. though we feel that they should be considered only as rough estimates rather Chan actual values and they reflect our inability to assess M (and £2.) during this epoch.," We put these values into the table, though we feel that they should be considered only as rough estimates rather than actual values and they reflect our inability to assess $\dot{M}$ (and $T_{wd}$ ) during this epoch." + This is because during the different phases of the outburst. additional components (such as the accretion disk ancl other obscurius malerial ejected. during outburst) contaminate aud veil (Longetal.20035). the spectrum of the white dwarL.," This is because during the different phases of the outburst, additional components (such as the accretion disk and other obscuring material ejected during outburst) contaminate and veil \citep{lon03b} + the spectrum of the white dwarf." + ILowever. during the cooling phase. one expects to see mainlv the white cwarf with little or no contribution from the disk and/or other (possibly masking) additional components.," However, during the cooling phase, one expects to 'see' mainly the white dwarf with little or no contribution from the disk and/or other (possibly masking) additional components." + During (his epoch the white dwarl is exposed and its temperature decreases., During this epoch the white dwarf is exposed and its temperature decreases. + Because of that. we decide to use for our modeling the values of the white dwarf temperature obtained during (he cooling phase only.," Because of that, we decide to use for our modeling the values of the white dwarf temperature obtained during the cooling phase only." + However. discrepancies of up to 5.000Ix exists between the temperature estimates of Sionetal...(2003a) and Longetal.(2003b).. for the observation carried out on day 50. when the white cwarl was revealed. but apparently also partially masked. or/amnd possibly with an accretion disk component.," However, discrepancies of up to 5,000K exists between the temperature estimates of \citet{sio03} and \citet{lon03b}, for the observation carried out on day 50, when the white dwarf was revealed, but apparently also partially masked or/and possibly with an accretion disk component." + The rest of the observations in the cooling phase are consisten within 200019. we elaborate here a little more on the temperature. using the three different approaches (denoted a. b and c in Table 2).," The rest of the observations in the cooling phase are consisten within 2,000K. we elaborate here a little more on the temperature, using the three different approaches (denoted a, b and c in Table 2)." + These were all observations using IIST/STIS., These were all observations using HST/STIS. + In (a) we mask the N V region of the spectrum when needed (together will less affected regions varving lrom spectrum to spectrum ) and use the latest version of the synthetic spectra generator codes TLUSTY and SYNSPEC (Ilubeny1988:IIubeny.οἱal.1994:IIubenyv&Lanz1995): in (b) the spectral fitting technique uses (he same masking regions for all the spectra together with (he new version ol the code: and (ο) is taken [rom Longetal...(2003b).. which uses a third. and different masking technique together with an earlier version of the code.," In (a) we mask the N V region of the spectrum when needed (together with less affected regions varying from spectrum to spectrum ) and use the latest version of the synthetic spectra generator codes TLUSTY and SYNSPEC \citep{hub88,hub94,hub95}; in (b) the spectral fitting technique uses the same masking regions for all the spectra together with the new version of the code; and (c) is taken from \citet{lon03b}, which uses a third and different masking technique together with an earlier version of the code." + In Table 2. we also list the [Iux integrated over the wavelength over (he entire spectral range of STIS. to the power 1/4. for all the epochs. which is proportional to the effective temperature (though each integrated flux can be relatively over- or under- estimated).," In Table 2, we also list the flux integrated over the wavelength over the entire spectral range of STIS, to the power 1/4, for all the epochs, which is proportional to the effective temperature (though each integrated flux can be relatively over- or under- estimated)." +" In Figure 2 we draw the temperatures listed in Table 2 as follows: Zi, is represented by stars: Z5 is represented by squares: 7; is represented by plus signs: and (the temperature estimated through the flix is represented by circles. where it has been arbitrarily scaled so as to fit (he last data point for 75."," In Figure 2 we draw the temperatures listed in Table 2 as follows: $T_a$ is represented by stars; $T_b$ is represented by squares; $T_c$ is represented by plus signs; and the temperature estimated through the flux is represented by circles, where it has been arbitrarily scaled so as to fit the last data point for $T_b$." + The two triangles represent the FUSE data points., The two triangles represent the FUSE data points. + The september 2001 datum has temperature estimates ranging from 221. 000Ix up to 232. 0001x..," The September 2001 datum has temperature estimates ranging from $\approx 27,000$ K up to $\approx 32,000$ K." +velocity dispersion during NBO erowth as a function of time.,velocity dispersion during KBO growth as a function of time. +" It is apparent from Figure 1. that the erowth and velocity evolution become selt£siuüblu once X—al241σ10'3 ""co. he. the shape of the size-cdistribution and velocity-distribution remains unchauged. while the uaxiuun KDO size coutiuues to erow."," It is apparent from Figure \ref{fig1} that the growth and velocity evolution become self-similar once $\Sigma \sim \alpha^{3/4} +\sigma \sim 10^{-3} \sigma$ , i.e. the shape of the size-distribution and velocity-distribution remains unchanged, while the maximum KBO size continues to grow." + The two upper panels of Figure 1. and Figure 2. show hat the size-distribution of large KBOs (R=50kin} iudeed follows a power-law with 4=Las predicted * our analytic treatment iu section ??.., The two upper panels of Figure \ref{fig1} and Figure \ref{fig2} show that the size-distribution of large KBOs $R \gtrsim 50~\rm{km}$ ) indeed follows a power-law with $q=4$ as predicted by our analytic treatment in section \ref{s1}. + This implies a roughly equal amount of mass per logarithmic mass in for hwge KDBOs., This implies a roughly equal amount of mass per logarithmic mass bin for large KBOs. +" Moreover. the middle panel of Figure Lo aud Fieure 2. show that the mass ratio of aree to sinall KBOs, i.e. S/o. found in our coagulation sinulation agrees very well with our analytic prediction hat δισ.αF07."," Moreover, the middle panel of Figure \ref{fig1} and Figure \ref{fig2} show that the mass ratio of large to small KBOs, i.e. $\Sigma/\sigma$, found in our coagulation simulation agrees very well with our analytic prediction that $\Sigma/\sigma \sim \alpha^{-3/4} \sim +10^{-3}$." + The sinmlatious confirm our analytic results and sugecst that the total mass in large objects that we see in the IXuiper belt is uot arbitrary it an outcome of the KBO evowth iud that it is roughly 10? of the initial planetesimeals mass., The simulations confirm our analytic results and suggest that the total mass in large objects that we see in the Kuiper belt is not arbitrary but an outcome of the KBO growth and that it is roughly $10^{-3}$ of the initial planetesimals mass. + This result is in exccllent aerecment with the actual observed mass iu aree WBOs aud formation from a MMSN type disk., This result is in excellent agreement with the actual observed mass in large KBOs and formation from a MMSN type disk. + Ow work therefore sugecsts that the Kuiper belt did ος contain two to three order of maguitude more mass in laree KBOs as has been proposed by some models (Weideuschilling2002:," Our work therefore suggests that the Kuiper belt did not contain two to three order of magnitude more mass in large KBOs as has been proposed by some models \citep{W02,TGM05}." +Tsiganisetal.2005).. Fiewre 3 shows a comparison of the velocity dispersion from) our coagulation simulation and our analytic results derived iu section ?7.., Figure \ref{fig3} shows a comparison of the velocity dispersion from our coagulation simulation and our analytic results derived in section \ref{s1}. + It. displavs very eood agreement between the velocity dispersion that we derived analytically for the various size regimes. aud the results from our coagulation simulation.," It displays very good agreement between the velocity dispersion that we derived analytically for the various size regimes, and the results from our coagulation simulation." + Figures 1-3. show that our analytic work captures the essential features of KDBO erowth aud that analytic theory and the nunerical coagulation results are iu exccllent aerecinent., Figures \ref{fig1}- \ref{fig3} show that our analytic work captures the essential features of KBO growth and that analytic theory and the numerical coagulation results are in excellent agreement. + We are able to successtiully explain the slope (¢= 1) aud amplitude (X~a the large KBO size distribution and the evolution of the velocity dispersion in the various velocity aud size roelnies., We are able to successfully explain the slope $q=4$ ) and amplitude $\Sigma \sim \alpha^{-3/4} \sigma$ ) of the large KBO size distribution and the evolution of the velocity dispersion in the various velocity and size regimes. + We performed an additional set of coagulation simulations with different initial conditious., We performed an additional set of coagulation simulations with different initial conditions. +" In the first set we started with most of the mass in simall 1-kui-sized objects. just as before. but added a second population of arger l0-kni-ized KDOs that contained 10.? of the total nass,"," In the first set we started with most of the mass in small 1-km-sized objects, just as before, but added a second population of larger 10-km-sized KBOs that contained $10^{-3}$ of the total mass." + All bodies were started with a velocity dispersion equal to 3 times their Tall velocity aud we followed heir growth aud velocity evolution., All bodies were started with a velocity dispersion equal to 3 times their Hill velocity and we followed their growth and velocity evolution. + Figure { shows he result of this erowth (points) and a comparison with the cumulative mass distribution at various times or KDBOs that erew from a sinele population of 1-kui sized planctesimals (lines)., Figure \ref{fig4} shows the result of this growth (points) and a comparison with the cumulative mass distribution at various times for KBOs that grew from a single population of 1-km sized planetesimals (lines). + The similarity between the wo distributions is striking., The similarity between the two distributions is striking. + The same power-law for he large IKBO size distribution emerges aud the mass ratio in lavee and simall KBOs becomes also the same in both simulations., The same power-law for the large KBO size distribution emerges and the mass ratio in large and small KBOs becomes also the same in both simulations. + These results highlight the power of he “two eroups approximation that we used to derive he analytic results in section ?? and validate our assertion that the growth of large KBOs develops towards a state where small and large bodies coutribute about equally to the erowth.," These results highlight the power of the `two groups approximation' that we used to derive the analytic results in section \ref{s1} + and validate our assertion that the growth of large KBOs develops towards a state where small and large bodies contribute about equally to the growth." + The overall growth timescale to, The overall growth timescale to +value will be related to (he fractional uncertainty in the lensing potential (0) and the square of its gradient. evaluated at the (wo SN image positions. (,"value will be related to the fractional uncertainty in the lensing potential $\psi (\vec{\theta})$ and the square of its gradient, evaluated at the two SN image positions. (" +This can be seen by substituting (2)) into (1)) to eliminate the source position j. ),This can be seen by substituting \ref{lens_eq}) ) into \ref{arrival_time}) ) to eliminate the source position $\vec{\beta}$ .) + Based on their nonparametric lensing inversions. AbdelSalam et ((1998a.b) report fractional uncertainties in the convergence & ol < in the strong lensing regions of interest in the clusters Abell 370 (using 6 svstems) and Abell 2218 (using 3 miultiple-imaged svstems and a number of singlv imaged. distorted arclets).," Based on their nonparametric lensing inversions, AbdelSalam et (1998a,b) report fractional uncertainties in the convergence $\kappa$ of $\le$ in the strong lensing regions of interest in the clusters Abell 370 (using 6 multiple-image systems) and Abell 2218 (using 3 multiple-imaged systems and a number of singly imaged, distorted arclets)." + They demonstrate clearly that the mass distribution becomes nore tightlv constrained as (he number of nearby lensed images increases., They demonstrate clearly that the mass distribution becomes more tightly constrained as the number of nearby lensed images increases. + The cramatic increase in multiply imaged sources that ACS and later instruments can uncover should ertainlv push the uncertainty in reconstructed cluster mass distributions well below in the inner cluster regions., The dramatic increase in multiply imaged sources that ACS and later instruments can uncover should certainly push the uncertainty in reconstructed cluster mass distributions well below in the inner cluster regions. + Furthermore. strong-lensing features constrain (he potential directly. and (he potential is an integral over the mass distribution.," Furthermore, strong-lensing features constrain the potential directly, and the potential is an integral over the mass distribution." + Ποιός we expect the uncertainty in (he cluster potential (the more relevant uncertaintv. for. Πρ determination) to be less (han the uncertainty in the reconstructed mass distribution., Hence we expect the uncertainty in the cluster potential (the more relevant uncertainty for $H_0$ determination) to be less than the uncertainty in the reconstructed mass distribution. + SN image positions (and perhaps flux ratios. if microlensing and differential extinction elfects are negligible) also provide local constraints on the potential. Farther improving the accuracy of a (ime-delay. Lf) value.," SN image positions (and perhaps flux ratios, if microlensing and differential extinction effects are negligible) also provide local constraints on the potential, further improving the accuracy of a time-delay $H_0$ value." + Taken together (hese considerations suggest that statistical //j errors on the order of a lew percent should be attainable if successful observations can be made., Taken together these considerations suggest that statistical $H_0$ errors on the order of a few percent should be attainable if successful observations can be made. + And as noted in the introduction. the lens model will not be subject to systematic errors associated with a elobal mass sheet degeneracy as long as mulliple-image constraints are available lor sources al different. recdshilts.," And as noted in the introduction, the lens model will not be subject to systematic errors associated with a global mass sheet degeneracy as long as multiple-image constraints are available for sources at different redshifts." + To estimate the error that could arise due to Gime-celay contributions from unmodelecd ealaxv-scale substructure. we note (hat time delavs scale as (he square of the characteristic deflection of the lens. which in (urn seales as the square of the lens velocity dispersion. as in (5)) and (6)).," To estimate the error that could arise due to time-delay contributions from unmodeled galaxy-scale substructure, we note that time delays scale as the square of the characteristic deflection of the lens, which in turn scales as the square of the lens velocity dispersion, as in \ref{theta_e}) ) and \ref{SIS_diff}) )." + With a time delay of 0.174. à cluster velocity dispersion of 1000km/s. and a perturbing galaxy. velocity dispersion of 200km/s. the induced fractional error in Lf) would be Thus substructure is unlikely to be a source of large error in the determination of Ly [rom an observed time delay.," With a time delay of $0.1 \tau_{\mathrm{cluster}}$ , a cluster velocity dispersion of $1000 \, \mathrm{km} / \mathrm{s}$, and a perturbing galaxy velocity dispersion of $200 \, \mathrm{km} / \mathrm{s}$, the induced fractional error in $H_0$ would be Thus substructure is unlikely to be a source of large error in the determination of $H_0$ from an observed time delay." + An appreciable substructure perturbation to the overall cluster potential would also affect. lensed SN image positions (which can (themselves constrain the potential). so it could not go entirely unmodeled by areconstruction Chat makes use of all lensing constraints.," An appreciable substructure perturbation to the overall cluster potential would also affect lensed SN image positions (which can themselves constrain the potential), so it could not go entirely unmodeled by areconstruction that makes use of all lensing constraints." +algorithm using only the PATTERN=0 events gives a position that is only 1.27 away.,algorithm using only the $PATTERN$ =0 events gives a position that is only 1.2” away. + One would immediately think of using nearby X-ray sources associated with well-known stars - sources which are less bright in X-rays (thus unaffected by pile-up) but still bright enough to get an accurate position in each dataset., One would immediately think of using nearby X-ray sources associated with well-known stars - sources which are less bright in X-rays (thus unaffected by pile-up) but still bright enough to get an accurate position in each dataset. + However. such nearby sources do not exist in the neighbourhood ofPuppis.," However, such nearby sources do not exist in the neighbourhood of." +. A perfect centroiding of annular regions is thus impossible and using always the Hippareos position for annular regions may alter the source’s properties. hence the results of the variability study that we wish to perform.," A perfect centroiding of annular regions is thus impossible and using always the Hipparcos position for annular regions may alter the source's properties, hence the results of the variability study that we wish to perform." +timescales than their main sequence counter-parts for which the RLOF is encountered. on longer (mnagnetic braking) timescales.,timescales than their main sequence counter-parts for which the RLOF is encountered on longer (magnetic braking) timescales. + The resulting rate of poteutial DDS SNe Ia varies substantially with tine., The resulting rate of potential DDS SNe Ia varies substantially with time. + At carly times (fX1d Cyr} the rates are very high ~0.01 4. and then they exaduallv decrease to reach ~0.0003 vr1 at late tines (f=X10 Gar).," At early times $t \lesssim 1$ Gyr) the rates are very high $\sim 0.01$ $^{-1}$, and then they gradually decrease to reach $\sim 0.0003$ $^{-1}$ at late times $t \gtrsim 10$ Gyr)." + The observed rates for elliptical galaxies are estimated at the level of πω~0.00184x0.0006 Lb per unit (019£2 3. of blue huninositv (Cappellaroal. 1999)., The observed rates for elliptical galaxies are estimated at the level of ${\cal R}_{\rm obs} \sim 0.0018 \pm 0.0006$ $^{-1}$ per unit $10^{10} L^B_\odot$ ) of blue luminosity \citep{CET99}. +. As the blue huninosity of elliptical ealaxies declines with time (after an carly star formation episode). the rates preseuted in the top paucl of Figure 2 should be corrected downwards at carly times. while at later times thev should be increased if our rates are to be compared with those of typical ellipticals.," As the blue luminosity of elliptical galaxies declines with time (after an early star formation episode), the rates presented in the top panel of Figure 2 should be corrected downwards at early times, while at later times they should be increased if our rates are to be compared with those of typical ellipticals." +" Obviously. the burst of star formation ou the order of6«1029AL. would produce a blue luninosity lavecr than 1019EZ, while 10-15 Cyr after the episode when stars more massive than 1M. have formed remnauts and are no longer contributing to the ealasxw’s light. the blue huuimositv is smaller than the normalising value."," Obviously, the burst of star formation on the order of $6 \times 10^{10} \msun$ would produce a blue luminosity larger than $10^{10} L^B_\odot$, while 10-15 Gyr after the episode when stars more massive than $\sim 1 \msun$ have formed remnants and are no longer contributing to the galaxy's light, the blue luminosity is smaller than the normalising value." + Since we do uot really know the distribution of age of the galaxies in the observed sample of cllipticals that were used iu the SN Ia rate estimate. we do not attempt to correct our svuthetic rates for the evolution ofblue huninosity aud we do not compare them directly to the observed rates of Cappollaroetal.(1999).," Since we do not really know the distribution of age of the galaxies in the observed sample of ellipticals that were used in the SN Ia rate estimate, we do not attempt to correct our synthetic rates for the evolution of blue luminosity and we do not compare them directly to the observed rates of \citet{CET99}." +" However, we note that the observed rate is consistent with our predicted rates for the DDS progenitor. while the predicted rates for other progenitors (SDS and AM CVn) seein to be significantly too low."," However, we note that the observed rate is consistent with our predicted rates for the DDS progenitor, while the predicted rates for other progenitors (SDS and AM CVn) seem to be significantly too low." + Iu Figure 2 (bottom panel). we show the SN In rates for the spiral galaxy model of Model 1.," In Figure 2 (bottom panel), we show the SN Ia rates for the spiral galaxy model of Model 1." + It is found that DDS rates of SNe Ia at the current epoch are 0.002 Ll, It is found that DDS rates of SNe Ia at the current epoch are $0.002$ $^{-1}$. + At first. the DDS rate increases with time (after the onset of star formation). then remains approsximatcly constant until the star formation stops leading to au overall decline in the rate.," At first, the DDS rate increases with time (after the onset of star formation), then remains approximately constant until the star formation stops leading to an overall decline in the rate." + This behavior reflects the specific shape of the delay time distribution for the DDS combined with the SER for our spiral galaxy model., This behavior reflects the specific shape of the delay time distribution for the DDS combined with the SFR for our spiral galaxy model. + The rates for SDS aud Κα progenitors are much smaller and at the level of ~104| 4., The rates for SDS and AM CVn progenitors are much smaller and at the level of $\sim 10^{-4}$ $^{-1}$. + SDS progenitors can eencrate SNe Ia long after star formation las ceased (long delay times). while AAD ΟΥ events disappear shortly after the star formation has stopped (short delay times).," SDS progenitors can generate SNe Ia long after star formation has ceased (long delay times), while AM CVn events disappear shortly after the star formation has stopped (short delay times)." + For comparison. over-plotted are empirical rates of SNe Ia. The rates were adopted from Cappollaroctal.(1999). for a Milkv. Way type spiral (Sbc-Sd) with a blue luninosity of 2«1079 L.. aud the rates are Row=0.0042350.002 SN Ta i," For comparison, over-plotted are empirical rates of SNe Ia. The rates were adopted from \citet{CET99} for a Milky Way type spiral (Sbc-Sd) with a blue luminosity of $2 \times 10^{10}$ $_{\odot}$, and the rates are ${\cal R}_{\rm obs}=0.004 \pm 0.002$ SN Ia $^{-1}$." + The DDS vate alone is consistent with the empirical rate of SNe Ia. The SDS and AM CVn SN In rates do not even come close to the enpirical rate. aud their addition to the DDS rate does uot sienificantly affect the overall rates at amy epoch.," The DDS rate alone is consistent with the empirical rate of SNe Ia. The SDS and AM CVn SN Ia rates do not even come close to the empirical rate, and their addition to the DDS rate does not significantly affect the overall rates at any epoch." + We note that our mass normalization which προς a constant star formation history for 10 Cor results in a SER at the level of GAL.vr. |., We note that our mass normalization which implies a constant star formation history for 10 Gyr results in a SFR at the level of $6 \mpy$ . + The elobal SFR in the AIW nav be somewhat lower: ~3.5AL.vr| (Cox2000:O'Shaughnessyetal. 2008).. and in that case the DDS rates are only mareinally consistent with the observed Cappellaroetal.(1999). rates.," The global SFR in the MW may be somewhat lower: $\sim 3.5 +\mpy$ \citep{Cox00,Osh08}, and in that case the DDS rates are only marginally consistent with the observed \citet{CET99} rates." + On the other haud. it has been suggested that the SFR of the MW. has been decreasing with tine. onlv reaching ~3.5M.vr bat the current epoch (Nelemansctal.2001.see2.2).," On the other hand, it has been suggested that the SFR of the MW has been decreasing with time, only reaching $\sim 3.5 \mpy$ at the current epoch \citep[][see sect. 2.2]{NYP01,NYP04}." +. I£ such an estimate had been used the average SER of the MW is found at the level of ~8M.vr+. ancl our results would scale up. being consistent with the DDS scenario as the major SN Ia contributor in ΑΝκο spiral ealaxies. as long as the Cappellaroetal.(1999) rates are ene used for Comparison.," If such an estimate had been used the average SFR of the MW is found at the level of $\sim 8 \mpy$, and our results would scale up, being consistent with the DDS scenario as the major SN Ia contributor in MW-like spiral galaxies, as long as the \citet{CET99} rates are being used for comparison." +2. There is a marked decrease. bv ucarly a actor of 2. in the total umuber of SNe Ia progenitors in our model with decreased CE removal efficiency.," There is a marked decrease, by nearly a factor of 2, in the total number of SNe Ia progenitors in our model with decreased CE removal efficiency." + The overall decrease is due to the fact that the most dominant xoteutial channel. the DDS. is ouly ~50 as cfficicut. since a larger fraction of binaries will merece in the cohbunon euvelope phase rather than surviving the CE o subsequently form a double white dwarf.," The overall decrease is due to the fact that the most dominant potential channel, the DDS, is only $\sim 50$ as efficient, since a larger fraction of binaries will merge in the common envelope phase rather than surviving the CE to subsequently form a double white dwarf." + Iu Figure 3 (top panel). SN Ta rates are shown for our Model 2 elliptical galaxy (degνA=0.5. instantaneous starburst at £20).," In Figure 3 (top panel), SN Ia rates are shown for our Model 2 elliptical galaxy $\alpha_{\rm CE} \times \lambda = 0.5$, instantaneous starburst at $t=0$ )." + DDS SNe Ia progenitors coutiuue to outnumber the SDS and AM. CVn progenitors. however there are some notable differences.," DDS SNe Ia progenitors continue to outnumber the SDS and AM CVn progenitors, however there are some notable differences." + For short delay times fol vr. the DDS rates are nearly a factor of 2 lower than they are for Model 1.," For short delay times $t \sim 1 $ Gyr, the DDS rates are nearly a factor of 2 lower than they are for Model 1." + Then at later times. the Model 2 DDS rates are a factor of ~3 below those of Model 1. reaching ~9&107 tat delay times of LO Cyr (vs. 3.10tyr | for Model 1).," Then at later times, the Model 2 DDS rates are a factor of $\sim 3$ below those of Model 1, reaching $\sim 9 \times 10^{-5}$ $^{-1}$ at delay times of 10 Gyr (vs. $3 \times 10^{-4}$ $^{-1}$ for Model 1)." + Despite the lower DDS rates of Model 2. potential progenitors are found at all delay times. as they are in Model 1.," Despite the lower DDS rates of Model 2, potential progenitors are found at all delay times, as they are in Model 1." + The CE efficiency ofModel l allows for DDS progenitors to be drawn from a wider range (goine to sinaller values) among the distribution of initial separations. where as progenitors with small initial separation in Model 2 are removed from the DDS population in mergers during the CE phase.," The CE efficiency of Model 1 allows for DDS progenitors to be drawn from a wider range (going to smaller values) among the distribution of initial separations, where as progenitors with small initial separation in Model 2 are removed from the DDS population in mergers during the CE phase." + The Model 2 SDS chaunel is more efficient (by a factor of 3) than the Model 1 SDS chaunel since post-CE WD | MS binaries are found ou closer orbits., The Model 2 SDS channel is more efficient (by a factor of 3) than the Model 1 SDS channel since post-CE WD + MS binaries are found on closer orbits. + One major difference between he elliptical galaxy Ia rates of Model 1 aud Model 2 is hat the SDS rates match those of the DDS rates for delay ues2.55.5 Gyr (0.0002 ve+).," One major difference between the elliptical galaxy Ia rates of Model 1 and Model 2 is that the SDS rates match those of the DDS rates for delay times $\sim +2.5 - 5.5$ Gyr $\sim 0.0002$ $^{-1}$ )." + Since the majority of the douors are evolved stars (giauts or sub-giauts). he delay involves two components: the maim sequence Ποιο of the donor (a few Cyr for a donor to become a lant) and the accretion timescale. over which the mhnary WD can increase its mass to the Chandrasekhar nass (10-100. Alyy).," Since the majority of the donors are evolved stars (giants or sub-giants), the delay involves two components: the main sequence lifetime of the donor (a few Gyr for a donor to become a giant) and the accretion timescale, over which the primary WD can increase its mass to the Chandrasekhar mass (10-100 Myr)." + Thus the main sequence Lifetime of the donor is what sets the delay times for the SDS DTD., Thus the main sequence lifetime of the donor is what sets the delay times for the SDS DTD. + Since the stars are found on closer orbits after he common envelope. RLOF is cucountered between tle WD and the nou-degeuerate conrpanion more often in Model 2 (typically when the donor is a sub-giaut).," Since the stars are found on closer orbits after the common envelope, RLOF is encountered between the WD and the non-degenerate companion more often in Model 2 (typically when the donor is a sub-giant)." + Rates of potential AM CVu SN Ta are lower than those of Model l for elliptical galaxies aud are at the level of ~0.0002 +., Rates of potential AM CVn SN Ia are lower than those of Model 1 for elliptical galaxies and are at the level of $\sim 0.0002$ $^{-1}$. + For the majority of the progenitors the delav times are very short. so as in Model 1. these type of events are expected only in voung host galaxies or in reeious witli ongoing star formation.," For the majority of the progenitors the delay times are very short, so as in Model 1, these type of events are expected only in young host galaxies or in regions with ongoing star formation." + ‘Fast’ AM. CVu progenitors are more rare nu Model 2 since these systems more readily morec iu one of the two common euvelope phases that lead to the formation of these progenitors., `Fast' AM CVn progenitors are more rare in Model 2 since these systems more readily merge in one of the two common envelope phases that lead to the formation of these progenitors. + Iu contrast to Model 1. there is a simall contribution of the slow’ AM CVn progenitors doug delay times) in Model 2.," In contrast to Model 1, there is a small contribution of the `slow' AM CVn progenitors (long delay times) in Model 2." + The AM CVnu channel is outummbcred by both the SDS aud DDS channels at all epochs in the Model 2 elliptical galaxy., The AM CVn channel is outnumbered by both the SDS and DDS channels at all epochs in the Model 2 elliptical galaxy. + Iu Figure 3 (bottom panel). SN Ta rates are shown for our Model 2spiral galaxy.," In Figure 3 (bottom panel), SN Ia rates are shown for our Model 2spiral galaxy." + It is found that DDS rate of, It is found that DDS rate of +reddening correction does affect the oxveen abundauce we derive by reducing the 11A3727 inteusity. but this effect has less mupact on the osvecn abuudauce than the uncertaiutv. in. the clectrou temperature since: there ix so little oxvecu iu the forma of I,"reddening correction does affect the oxygen abundance we derive by reducing the $\lambda$ 3727 intensity, but this effect has less impact on the oxygen abundance than the uncertainty in the electron temperature since there is so little oxygen in the form of ${}^{+}$." + Third. forciug IOIS/I(OLJ)=0.17 via a reddeningO correction. even if negative. accounts for any οπου in the scusitivity calibration that ασ otherwise systematically affect the ΟΠΗ lines aud. the subsequent oxygen. abundances.," Third, forcing $I(\mathrm{H}\gamma)/I(\mathrm{H}\beta)=0.47$ via a reddening correction, even if negative, accounts for any errors in the sensitivity calibration that might otherwise systematically affect the ] lines and the subsequent oxygen abundances." + Fourth. on average. our IHa- aud IH5-based reddeuiugs agree.," Fourth, on average, our $\alpha$ - and $\gamma$ -based reddenings agree." +" The mean o-based reddening for all objects (both and \32))is LEBV)=O1840.0 πας, while the moan I2-based reddening for all ofthe planetary nebulae in the bulee of is LOBV)=0.18£0.08 nunag. if neeative reddening values are iucluded. or E(BV)=0.35+0.06 nanag. if negative reddening values are κο to zero (the uncertainties are the standard errors in the means)."," The mean $\alpha$ -based reddening for all objects (both and ) is $E(B-V)=0.18\pm 0.04$ mag, while the mean $\gamma$ -based reddening for all of the planetary nebulae in the bulge of is $E(B-V)=0.18\pm 0.08$ mag, if negative reddening values are included, or $E(B-V)=0.25\pm 0.06$ mag, if negative reddening values are set to zero (the uncertainties are the standard errors in the means)." + Thus. the reddenines conued from Πα aud II are simular.," Thus, the reddenings computed from $\alpha$ and $\gamma$ are similar." +" For comparison. t16 foreeround reddening to ABLis E(B.V)—0,09340.0an nae (mean of MeClur Racine 1969.. an den Dergh 1969.. and Burstein IIeiles 198 L9)."," For comparison, the foreground reddening to is $E(B-V)=0.093\pm 0.009$ mag (mean of McClure Racine \cite{McClureRacine1969}, van den Bergh \cite{vandenBergh1969}, and Burstein Heiles \cite{BursteinHeiles1984}) )." + It is not surprising hat the mean reddent for the planctary neπας is nuuag ercater than the foreground value. for plauetary neπι]αος suffer additional reddening due tointernal dus a-d dust within axdAD32.," It is not surprising that the mean reddening for the planetary nebulae is mag greater than the foreground value, for planetary nebulae suffer additional reddening due tointernal dust and dust within and." +. Cousequenthy. we have chosen to correct for reddeniusC» even when E(BVW) is negative.," Consequently, we have chosen to correct for reddening"" even when $E(B-V)$ is negative." +"Oo Tables Land 5 prescut the electron temperatures aud the oxveen abundances for the planetary nebulae in aud in the bulee ofM31.. respectively,"," Tables \ref{table6} and \ref{table7} present the electron temperatures and the oxygen abundances for the planetary nebulae in and in the bulge of, respectively." + We only observed two lonization stages of oxygen. O! and Of.," We only observed two ionization stages of oxygen, ${}^{+}$ and ${}^{++}$." + We acconutefor unseen sages dn our oxveen abindauce calculationS using the ionization correctkn factrs (ICE) coluputcc according to the prescription of Ixiugsbureh Barlow (] 991). which eumplovs tjo line intensities of À 1686 aid 5876 to ¢Orrect ‘Ok uuseen jonmzatiol stages of ONVECL.," We accountedfor unseen stages in our oxygen abundance calculations using the ionization correction factors (ICF) computed according to the prescription of Kingsburgh Barlow \cite{KingsburghBarlow1994}) ), which employs the line intensities of $\lambda$ 4686 and $\lambda$ 5876 to correct for unseen ionization stages of oxygen." + Further details mmy be fouud il Stasinsska e al. (1998))., Further details may be found in Stasińsska et al. \cite{Stasinskaetal1998}) ). + Tables aud 5 presen two oxvecu abundance calculations., Tables \ref{table6} and \ref{table7} present two oxygen abundance calculations. + The abuwlances 1l coluniu 35H are snplv the siu of t1ο QO! aud MI lonic abundances., The abundances in column 3 are simply the sum of the ${}^{+}$ and ${}^{++}$ ionic abundances. + The abuudauces in colanu lare hose from COmun 3 corrected for the ICE., The abundances in column 4 are those from column 3 corrected for the ICF. + The ICF is normally sia1 because Tel ALGS6 is weak., The ICF is normally small because $\lambda$ 4686 is weak. + The oxvecn abuidances 1 Cyuma Lwill be adopted im fuure work., The oxygen abundances in column 4 will be adopted in future work. + Iu calculating the oxvecn amudances. we assumed a1 electron density of 00002.7 in all cases.," In calculating the oxygen abundances, we assumed an electron density of $4000\,{\rm cm}^{-3}$ in all cases." + Wih electro1 desities of loni5m and 210502. 7. the oxvgen abundance changes by a maxima of 1.02 ddex: aud 0.07 4ος. respectively. for the planetary nebulae in ALL. aud by a," With electron densities of $1\,{\rm cm}^{-3}$ and $2\,10^{4}{\rm cm}^{-3}$ , the oxygen abundance changes by a maximum of $-0.02$ dex and $+0.07$ dex, respectively, for the planetary nebulae in , and by a" +however. that Imanishi Ueno (2000) argued in favour of a