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. The scenario envisaged by Guetta ancl Granot is unique in (hat it predicts TeV neutrinos from nuclear 'Zrocesses. | The scenario envisaged by Guetta and Granot is unique in that it predicts TeV neutrinos from nuclear processes. |
This is because of the presence in their model of protons accelerated to very. hieh Ijergies in a pulsar wind bubble. | This is because of the presence in their model of protons accelerated to very high energies in a pulsar wind bubble. |
The multiGeV collapsar neutrino signal we study here is related (0 the mechanism eiving rise to the variation in Lorentz factor of the outgoing jet. | The multi-GeV collapsar neutrino signal we study here is related to the mechanism giving rise to the variation in Lorentz factor of the outgoing jet. |
As with anv GRD model invoking internal shocks. (he jet variability must be substantial in order to explain observed temporal variability in GRBs. as well as to allow ellicient conversion of kinetic energy. (o electromagnetic energy (see Piran(1999). for a general overview). | As with any GRB model invoking internal shocks, the jet variability must be substantial in order to explain observed temporal variability in GRBs, as well as to allow efficient conversion of kinetic energy to electromagnetic energy (see \cite{pir99} for a general overview). |
In most GRD models. the jet variabilitv is related to a variability in the mass ablation rate or luminosity of the central source driving the outflow. | In most GRB models, the jet variability is related to a variability in the mass ablation rate or luminosity of the central source driving the outflow. |
By contrast. the jet variability in thecollapsar can arise from instabilities along the jet-stellar wall interface (Zhang.Woosley.&\laeFaclven2003). | By contrast, the jet variability in thecollapsar can arise from instabilities along the jet-stellar wall interface \citep{zha02}. |
. These instabilities operate at radi large compared to the size of the central black hole. | These instabilities operate at radii large compared to the size of the central black hole. |
Consequently. (he instabilities act to slow a jet which has already hac ample time to reach ultra-relativistic velocities. | Consequently, the instabilities act to slow a jet which has already had ample time to reach ultra-relativistic velocities. |
We point out that when free neutrons are present in (he jet. (his slowing is not elastic and is accompanied by a sizeable neutrino flux. | We point out that when free neutrons are present in the jet, this slowing is not elastic and is accompanied by a sizeable neutrino flux. |
In (he next section we discuss details of how the slowing of the neutrons in the jet leads to the conversion of jet energv to potentially observable neutrinos. | In the next section we discuss details of how the slowing of the neutrons in the jet leads to the conversion of jet energy to potentially observable neutrinos. |
Estimates of neutrino detection rates and energies are presented. | Estimates of neutrino detection rates and energies are presented. |
We will show that the neutrino signal of the slowing of the jet can originate Ilrom radii some (wo orders of magnitude smaller (han the signal of processes discussed by Mészáros&Rees (2000)... | We will show that the neutrino signal of the slowing of the jet can originate from radii some two orders of magnitude smaller than the signal of processes discussed by \cite{mes00}. . |
Also. the neutrino production we | Also, the neutrino production we |
Iz. | Hz. |
For the high state data a single power-law was fitted in the eutire frequency rauge. | For the high state data a single power-law was fitted in the entire frequency range. |
The distinguishiug features of the spectral states are 1) a flat spectrum below 1 Iz for the low-hard state with the power-law index ranging frou -0.5 to -0.2 compared to an index ~ -1.3 in the high-soft state. 2) presence of a narrow ος 10 Iz QPO feature in the low-hard state aud 3) higher variability (154 )) iu the 1 10 Tz rauge in the low-hard state compared to the low variability (<5% }) im the same frequeucy ranee for the ligh-soft state. | The distinguishing features of the spectral states are 1) a flat spectrum below 1 Hz for the low-hard state with the power-law index ranging from -0.5 to -0.2 compared to an index $\sim$ -1.3 in the high-soft state, 2) presence of a narrow 0.5 – 10 Hz QPO feature in the low-hard state and 3) higher variability $\sim$ ) in the 1 $-$ 10 Hz range in the low-hard state compared to the low variability $<$ ) in the same frequency range for the high-soft state. |
The other distinguishing characteristics of the iuteusity states of Calactic black hole. caudidate sources ds the οποιον spectra. | The other distinguishing characteristics of the intensity states of Galactic black hole candidate sources is the energy spectra. |
The low state is dominated by a thenual-Compton spectrum (which can be approximated to a power-law at lower energies) alone with a blackbody Cluission component. which egencrally is modeled as a disk blackbodyw cussion. | The low state is dominated by a thermal-Compton spectrum (which can be approximated to a power-law at lower energies) along with a blackbody emission component, which generally is modeled as a disk blackbody emission. |
The disk blackhody cussion Πιοαπο nni luteusitv and dominates the spectrum in the high state. | The disk blackbody emission increases in intensity and dominates the spectrum in the high state. |
We have eeucraed the 128 chainel euerev spectra from the Standard 2 mode of the PCA for cach of the above observations. | We have generated the 128 channel energy spectra from the Standard 2 mode of the PCA for each of the above observations. |
Staidird procedires for data selection. background estination and response matrix eeneration have been applic. | Standard procedures for data selection, background estimation and response matrix generation have been applied. |
To avoid the extra systematic errors iu the response matrix of PCA. we have restricted our analysis to the energy range of 3 26 keV. Data from all the PCUs are added together. | To avoid the extra systematic errors in the response matrix of PCA, we have restricted our analysis to the energy range of 3 $-$ 26 keV. Data from all the PCUs are added together. |
We have fitted the energv spectrun of the source using a model consisting of dixk-blackbodyw and power-law with absorption bx intervening cold material parameterized as cequivaleut Uvdrogen cobluun deusitv. Ny. | We have fitted the energy spectrum of the source using a model consisting of disk-blackbody and power-law with absorption by intervening cold material parameterized as equivalent Hydrogen column density, $_H$. |
The value of Ny has been kept fixed at 6 \< 1072 7. | The value of $_H$ has been kept fixed at 6 $\times$ $^{22}$ $^{-2}$. |
We have included a Ciaussian line near the expected Is,, emission frou iron and absorption edge due to iron. | We have included a Gaussian line near the expected $_\alpha$ emission from iron and absorption edge due to iron. |
These features help to mimic the reflection spectrun usually found iu other Calactic black hole candidate sources like Cwveguus N-1 (Cüerliuski et al. | These features help to mimic the reflection spectrum usually found in other Galactic black hole candidate sources like Cygnus X-1 (Gierlinski et al. |
1997). | 1997). |
Systematic errors of lave heen added to the data. | Systematic errors of have been added to the data. |
We obtain reduced 47 values iu the ranee of 1 3. | We obtain reduced $\chi^2$ values in the range of 1 – 3. |
The untoldecd spectrum for the two spectral states are shown in Figure 1. | The unfolded spectrum for the two spectral states are shown in Figure 4. |
The individual model components are shown separately. | The individual model components are shown separately. |
Scales for the axes are kept the same for all the plots. | Scales for the axes are kept the same for all the plots. |
The derived parameters are given iu Table 3. | The derived parameters are given in Table 3. |
The quoted errors are nominal coufiderce levels obtained by the condition of 2,5, | . | The quoted errors are nominal confidence levels obtained by the condition of $\chi_{min}^2$ + 2.7. |
The inner disk radius is caleulated using au immcliuation augle of 70? for the accretion disk aud a distance of 10 kpc to the source (see Muno et al. | The inner disk radius is calculated using an inclination angle of $^\circ$ for the accretion disk and a distance of 10 kpc to the source (see Muno et al. |
1999). | 1999). |
The iron Bue was found to have an equivalent width of 100 eV. We can ideutifv the following distinguishing features in the energv spectrum in the two spectral states of the source: 1) iu the low-hard state the disk blackbody component has lower temperature (<1 keV) aud larger iuner disk radius (2505 kin) compared to the high-soft state. which has the iuner disk temperature of ~2 keV aud immer disk radius of —20 kin. | The iron line was found to have an equivalent width of $\sim$ 100 eV. We can identify the following distinguishing features in the energy spectrum in the two spectral states of the source: 1) in the low-hard state the disk blackbody component has lower temperature $<$ 1 keV) and larger inner disk radius $>$ 50 km) compared to the high-soft state, which has the inner disk temperature of $\sim$ 2 keV and inner disk radius of $\sim$ 20 km. |
2) the disk blackbody component las a 3 26 keV flux of —10 9 ere. 72 s31 (<10% of the tota flux) whereas the disς blackbody flux iucreases bv a factor of more than 30 iu th high-sott state. | 2) the disk blackbody component has a 3 $-$ 26 keV flux of $\sim$ $^{-9}$ erg $^{-2}$ $^{-1}$ $<$ of the total flux) whereas the disk blackbody flux increases by a factor of more than 30 in th high-soft state. |
In fact it becomes the predominant component with >65% of the observed flux beimeOo in this couponeut. | In fact it becomes the predominant component with $>$ of the observed flux being in this component. |
3) The power-law index becomes uoticcably steeper iu the ligh-soft state. | 3) The power-law index becomes noticeably steeper in the high-soft state. |
Now we show that during the iegular bursts observed ou 1997 June 18 the source GRS 19151105 made rapid intensity transitions between two levels aud the X-rav enüssion properties iu these two levels are identical to the spectral states of the source. | Now we show that during the irregular bursts observed on 1997 June 18 the source GRS 1915+105 made rapid intensity transitions between two levels and the X-ray emission properties in these two levels are identical to the spectral states of the source. |
Belloui et al. ( | Belloni et al. ( |
19971) have shown that durius the regular bursts observed on 1997 June Ls the source exhibited a repeating pattern of inteusitv states characterized by well defined spectral states. | 1997b) have shown that during the irregular bursts observed on 1997 June 18 the source exhibited a repeating pattern of intensity states characterized by well defined spectral states. |
We have selected one ivreeular burst frou this observation for the timime aud spectral analysis. | We have selected one irregular burst from this observation for the timing and spectral analysis. |
The power deusitv spectra (PDS) were separately obtained | The power density spectrum (PDS) were separately obtained |
reasons in advance: the Na and especially the O abundances in giants are difficult to measure. and indeed we often see only upper limits for [O/Fe] or C. O. and Na are not directly comparable. because they are not produced exactly at the same temperature: C and N are altered within the CN bi-cycle. O is depleted in the complete CNO cycle and Na is produced in the NeNa cycle. each dominatingσι at progressively higher temperatures: i1 this second hypothesis much care should be taken with the comparison of Na-O and C-N anti-correlations with each other. | reasons in advance: the Na and especially the O abundances in giants are difficult to measure, and indeed we often see only upper limits for [O/Fe] or C, N, O, and Na are not directly comparable, because they are not produced exactly at the same temperature: C and N are altered within the CN bi-cycle, O is depleted in the complete CNO cycle and Na is produced in the NeNa cycle, each dominating at progressively higher temperatures; in this second hypothesis much care should be taken with the comparison of Na-O and C-N anti-correlations with each other. |
Finally we note that if such a C-N bimodality ts present. it should be easier to see in MS stars than in RGB stars. where it could be attenuated by internal mixing (Sect. 1)). | Finally we note that if such a C-N bimodality is present, it should be easier to see in MS stars than in RGB stars, where it could be attenuated by internal mixing (Sect. \ref{sec-intro}) ). |
We studied the presence of bimodal distributions for the CN and CH rectified band strengths. by means of histograms (Fig. 6)) | We studied the presence of bimodal distributions for the CN and CH rectified band strengths by means of histograms (Fig. \ref{fig_hist}) ) |
and — more importantly — by plotting our measurements in the CH-CN plane (Fig. 7.. | and – more importantly – by plotting our measurements in the CH-CN plane (Fig. \ref{fig_chcn}, |
see also next Sect.). | see also next Sect.). |
We consider a distribution to be "bimodal" when the centroids of the CN-strong (CH-weak) and CN-weak stars are clearly in the CH-CN plane and — as a supporting evidence — when the histograms also show signs of bimodality. | We consider a distribution to be “bimodal" when the centroids of the CN-strong (CH-weak) and CN-weak stars are clearly in the CH-CN plane and – as a supporting evidence – when the histograms also show signs of bimodality. |
Detailed studies of the statistical significance of such bimodalities are deferred to a following paper. dealing with the [C/Fe] and [N/Fe] distributions. | Detailed studies of the statistical significance of such bimodalities are deferred to a following paper, dealing with the [C/Fe] and [N/Fe] distributions. |
First. we note that all the most metal-rich clusters (Pal 12. 47 Tuc. NGC 6352. and NGC 5927). with the only exception of NGC 6388 (which has low S/N and few stars). show a clear bimodalitynot only n 0S3839(CN). but also in 0CH4300. | First, we note that all the most metal-rich clusters (Pal 12, 47 Tuc, NGC 6352, and NGC 5927), with the only exception of NGC 6388 (which has low S/N and few stars), show a clear bimodalitynot only in $\delta$ S3839(CN), but also in $\delta$ CH4300. |
Among the metal-poor clusters. the situation is less clear. | Among the metal-poor clusters, the situation is less clear. |
There is a bimodality only in M 19. even if it is the most metal-poor of the sample ael has a relatively low S/N ratio. | There is a bimodality only in M 15, even if it is the most metal-poor of the sample and has a relatively low S/N ratio. |
For the remaining clusters we only see vague hints of possible bimodalities (see Sect. | For the remaining clusters we only see vague hints of possible bimodalities (see Sect. |
5. for more details) in Fig. 6.. | \ref{sec-sample} for more details) in Fig. \ref{fig_hist}, |
but not in Fig. | but not in Fig. |
7 or 8.. | \ref{fig_chcn} or\ref{fig_zoom}. . |
We will disctSS ndividual clusters in Sect. 5.. | We will discuss individual clusters in Sect. \ref{sec-sample}. . |
required a more uncertain extrapolation beyond the limits of the theoretical models. | required a more uncertain extrapolation beyond the limits of the theoretical models. |
Histograms with the mass and age distribution for the 680 stars measured in this way are shown in reffig7aa and 7bb. respectively. | Histograms with the mass and age distribution for the 680 stars measured in this way are shown in \\ref{fig7}a a and \ref{fig7}b b, respectively. |
As reffig6 already implied. the histogram in reffig7bb shows that there is a remarkable paucity of stars with ages around MMvr. | As \\ref{fig6} already implied, the histogram in \\ref{fig7}b b shows that there is a remarkable paucity of stars with ages around Myr. |
Besides the Ha luminosity L(Hà0) already discussed in the previous section. the other physical parameters of interest in this work are the accretion luminosity {ως and the mass accretion rate My... that we derived following the procedure described in II. From an analysis of the literature values of La and L(Ho) measurements as recently summarised by Dahm (2008) for PMS stars in. the Taurus-Auriga association. that paper concluded that the ratio {ο(Πα) can be assumed to be constant. | Besides the $\alpha$ luminosity $L(H\alpha)$ already discussed in the previous section, the other physical parameters of interest in this work are the accretion luminosity $L_{\rm acc}$ and the mass accretion rate $\dot M_{\rm acc}$, that we derived following the procedure described in I. From an analysis of the literature values of $L_{\rm acc}$ and $L(H\alpha)$ measurements as recently summarised by Dahm (2008) for PMS stars in the Taurus–Auriga association, that paper concluded that the ratio $L_{\rm acc} / L(H\alpha)$ can be assumed to be constant. |
The proportionality constant. obtained from an elementary fit to the data in the compilation of Dahm (2008). is where the large uncertainty on the proportionality factor arises because the two quantities were determined from non-simultaneous observations. | The proportionality constant, obtained from an elementary fit to the data in the compilation of Dahm (2008), is: where the large uncertainty on the proportionality factor arises because the two quantities were determined from non-simultaneous observations. |
Considering that the intensity of the Πα line from PMS sequence stars is known to varyby about and by as much as a factor of 2 — 3 in a few days (e.g. Fernandez et al. | Considering that the intensity of the $\alpha$ line from PMS sequence stars is known to vary by about and by as much as a factor of 2 – 3 in a few days (e.g. Fernandez et al. |
1995; Smith et al. | 1995; Smith et al. |
1999; Alencar et al. | 1999; Alencar et al. |
2001). the large spread can be understood. | 2001), the large spread can be understood. |
However. as we will show in refevolu.. the small dispersion of the mass accretion rate with stellar age suggests an uncertainty smaller than 0.25. | However, as we will show in \\ref{evolu}, the small dispersion of the mass accretion rate with stellar age suggests an uncertainty smaller than $0.25$. |
The mass accretion rate is related to £4 via the free-fall equation. that links the luminosity released in theimpact of the accretion flow with the rate of mass accretion Macc. according to the relationship: where G is the gravitational constant. M. and R. the mass and photospheric radius of the star as derived above and Ri, the inner radius of the accretion disc. | The mass accretion rate is related to $L_{\rm acc}$ via the free-fall equation, that links the luminosity released in theimpact of the accretion flow with the rate of mass accretion $\dot M_{\rm acc}$, according to the relationship: where $G$ is the gravitational constant, $M_*$ and $R_*$ the mass and photospheric radius of the star as derived above and $R_{\rm in}$ the inner radius of the accretion disc. |
The value of Rj, 1s rather uncertain and depends on how exactly the accretion dise is coupled with the magnetic field of the star. | The value of $R_{\rm in}$ is rather uncertain and depends on how exactly the accretion disc is coupled with the magnetic field of the star. |
Following Gullbring et al. ( | Following Gullbring et al. ( |
1998). we adopt Ai,25R. for all PMS objects. | 1998), we adopt $R_{\rm in} = 5\,R_*$ for all PMS objects. |
By substituting refeqd in refeq5.. the mass accretion rate Mace. in units of can be directly expressed as a function of {(Πα) The histograms showing the distribution of LC0) and Mace as derived with refeq6 for the 680 bona-fide PMS stars with well defined masses are shown in reffig7ee and 7dd. respectively. | By substituting \\ref{eq4} in \\ref{eq5}, the mass accretion rate $\dot M_{\rm acc}$, in units of $^{-1}$, can be directly expressed as a function of $L(H\alpha)$ : The histograms showing the distribution of $L(H\alpha)$ and $\dot M_{\rm
acc}$ as derived with \\ref{eq6} for the 680 bona-fide PMS stars with well defined masses are shown in \\ref{fig7}c c and \ref{fig7}d d, respectively. |
The median values of L(Ha) and Mace are. respectively. 4.6«10?! eerg/s (or 0.0I2LL ..) and3.9«10"5 yr!. | The median values of $L(H\alpha)$ and $\dot
M_{\rm acc}$ are, respectively, $4.6 \times 10^{31}$ erg/s (or $0.012$ $_\odot$ ) and $3.9 \times 10^{-8}$ $^{-1}$. |
The latter value is about of a population of 133 PMS stars that we detected in the field of 11987A (see II) and this is fully consistent with the much younger average age of the PMS stars in 3346. | The latter value is about of a population of 133 PMS stars that we detected in the field of 1987A (see I) and this is fully consistent with the much younger average age of the PMS stars in 346. |
PaperIL provides an extensive discussion of the statistical and systematic. uncertainties involved in determining Mace with this method and we refer the reader to that work for more details. | I provides an extensive discussion of the statistical and systematic uncertainties involved in determining $\dot M_{\rm acc}$ with this method and we refer the reader to that work for more details. |
The typical combined statistical uncertainty on M for the 3346 field is L(Ha) of the photometry (we note that this uncertainty is lower than for the study of the field around 11987A in II because of the higher quality of the ACS photometry). | The typical combined statistical uncertainty on $\dot M_{\rm acc}$ for the 346 field is $L(H\alpha)$ of the photometry (we note that this uncertainty is lower than for the study of the field around 1987A in I because of the higher quality of the ACS photometry). |
As for systematic effects. those on R... M. and reddening add up to uncertain relationship. linking (Πα) to the.aceretion lummosity {εως that plays a dominant role. | As for systematic effects, those on $R_*$ , $M_*$ and reddening add up to uncertain relationship linking $L(H\alpha)$ to theaccretion luminosity $L_{\rm acc}$ that plays a dominant role. |
We stress here that our method adopts the calibration provided by Dahm (2008). which is based on a spectroscopic study of Galactic PMS | We stress here that our method adopts the calibration provided by Dahm (2008), which is based on a spectroscopic study of Galactic PMS |
wide range of properties, with best-fit stellar masses ranging from 10—10!"A.. ages from ~à—1000 Myr and metallicities from 0.005—1Z. (e.g. Charyetal.2005;LaiFinkelsteinetal. 2009)). | wide range of properties, with best-fit stellar masses ranging from $\sim 10^8 -10^{10} \: M_\odot$, ages from $\sim 3 - 1000$ Myr and metallicities from $0.005 - 1 \: Z_\odot$ (e.g. \citealp{chary05,lai07,finkelstein09a}) ). |
Evidence has also been found for significant amounts of dust in LAEs, with UV extinctions of τν~0.5—5 mag. ( | Evidence has also been found for significant amounts of dust in LAEs, with UV extinctions of $A_{\rm UV} \sim 0.5 - 5$ mag. ( |
e.g. Charyetal.2005:Finkelstein 2009)). | e.g. \citealp{chary05,finkelstein09a}) ). |
Charyetal.(2005) suggest that the observed decrease in the the cosmic star-formation rate at iVcC ~[e(f. Stanwayetal. 2004)) might be explained by the presence of dust in the 2>6 LAE population. | \citet{chary05} suggest that the observed decrease in the the cosmic star-formation rate at $z > 6$ (e.g. \citealp{stanway04}) ) might be explained by the presence of dust in the $z > 6$ LAE population. |
Typical star-formation rates (SFRs) in LAEs, as inferred from their UV continua, are low, a few tens of solar masses per year (c.g. Taniguchietal. 2005)); these are 2—5 times larger than the SFRs derived from the resonantly-scattered Lyman-a line. | Typical star-formation rates (SFRs) in LAEs, as inferred from their UV continua, are low, a few tens of solar masses per year (e.g. \citealp{taniguchi05}) ); these are $2-5$ times larger than the SFRs derived from the resonantly-scattered $\alpha$ line. |
While this difference could arise from dust obscuration effects, aabsorption by the damping wing of the Gunn-Petersen trough is likely to also contribute towards reducing the strength of the Lyman-a line for LAEs at z>6.5 (Haiman2002). | While this difference could arise from dust obscuration effects, absorption by the damping wing of the Gunn-Petersen trough is likely to also contribute towards reducing the strength of the $\alpha$ line for LAEs at $z > 6.5$ \citep{haiman02}. |
. We note that the tentative detection of the Ha line in a single z~6.56 LAE (Charyetal.2005) yielded a significantly higher SFR estimate (140M. yr |) than that obtained even from the rest-frame UV continuum (~9M. yr +): this emphasizes the possibility that the SFRs in other LAEs might have been under-estimated due to dust extinction. | We note that the tentative detection of the $\alpha$ line in a single $z \sim 6.56$ LAE \citep{chary05} yielded a significantly higher SFR estimate $140 \: M_\odot$ $^{-1}$ ) than that obtained even from the rest-frame UV continuum $\sim 9\: M_\odot$ $^{-1}$ ); this emphasizes the possibility that the SFRs in other LAEs might have been under-estimated due to dust extinction. |
The detectability of high--: galaxies like the LAEs relies on their undergoing an elevated level of star-formation activity, which naturally requires fuel in the form of molecular gas. | The detectability of $z$ galaxies like the LAEs relies on their undergoing an elevated level of star-formation activity, which naturally requires fuel in the form of molecular gas. |
Such gas is most effectively studied through observations of redshifted CO emission lines (e.g. VandenBout 2005)). | Such gas is most effectively studied through observations of redshifted CO emission lines (e.g. \citealp{solomon05}) ). |
The luminosity in the CCO lines can be used to estimate the total molecular gas mass fueling the star-formation activity, while the CO line widths provide a measure of the dynamical mass of the galaxy. | The luminosity in the CO lines can be used to estimate the total molecular gas mass fueling the star-formation activity, while the CO line widths provide a measure of the dynamical mass of the galaxy. |
Studies of molecular gas at high redshifts, >4, have so far focused on the most massive, far-infrared-Iuminous systems, the sub-mm galaxies or quasars (c.g. 2003)), and no information is available in the literature on the molecular gas content of "normal" star-forming galaxies, such as the LAEs. | Studies of molecular gas at high redshifts, $z > 4$, have so far focused on the most massive, far-infrared-luminous systems, the sub-mm galaxies or quasars (e.g. \citealp{schinnerer08,walter03}) ), and no information is available in the literature on the molecular gas content of “normal” star-forming galaxies, such as the LAEs. |
Fie. | Fig. |
3 shows a radia plot of an oserved (sky subtracted’) stellar xofile (square sviubols). obtained from a dieiINC plate. plotted as arbitrary inteusitv versus 5.25 are second pixels. | \ref{profile_comparison} shows a radial plot of an observed (`sky subtracted') stellar profile (square symbols), obtained from a digitised plate, plotted as arbitrary intensity versus 5.25 arc second pixels. |
The observed profile is saturated out ο abot το Or 7 pixels in radius. | The observed profile is saturated out to about 6 or 7 pixels in radius. |
Also potted is part oftje poly1onial fit to the Ning LO.file from Fig. 1)). | Also plotted is part of the polynomial fit to the King profile (from Fig. \ref{King}) ), |
wit Lasatyration limit 1iuposed so tiat the radius of satura1011 Is jxxels to approximately match the olserved profile. aud then scaed verically so that the saturaion iutcusitics late1. | with a saturation limit imposed so that the radius of saturation is 6 pixels to approximately match the observed profile, and then scaled vertically so that the saturation intensities match. |
Tje agreenient is reasolaable alt1O1eh the observed profie falls nore rapidly than the svithetie profile at larecr radii. | The agreement is reasonable although the observed profile falls more rapidly than the synthetic profile at larger radii. |
Tιο difference. at large radius is. at least in part. probably à coISCquUeice of r polvnonual ft (see Fie. 1)). | The difference at large radius is, at least in part, probably a consequence of our polynomial fit (see Fig. \ref{King}) ), |
which apICALS Sightly high near LOO arc sec (~20 pixels). | which appears slightly high near 100 arc sec $\sim$ 20 pixels). |
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