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In τοςAS owe discuss the three structural parameters of the CAS system. Concentration (C). symmetry. CX)and Clumpiness (5). | In \\ref{CAS} we discuss the three structural parameters of the CAS system, Concentration (C), Asymmetry (A)and Clumpiness (S). |
We then compare the simulations with the empirical ? sample of galaxies in refC'omparison.. | We then compare the simulations with the empirical \citet{Frei1996} sample of galaxies in \\ref{Comparison}. |
The results are discussed in refDiscuss and conclusions drawn in re[Concl.. | The results are discussed in \\ref{Discuss} and conclusions drawn in \\ref{Concl}. |
‘To compare with the morphological parameters. inferred from observed galaxies. we used three samples of simulated galaxies. | To compare with the morphological parameters inferred from observed galaxies, we used three samples of simulated galaxies. |
Each sample was generated using the gravitational N-Bodyv | smoothecl particle hiycrodyvnamies (SPLL) code (2).. | Each sample was generated using the gravitational N-Body + smoothed particle hydrodynamics (SPH) code \citep{Wadsley2004}. |
We brielly outline the simulations here. but refer the reader to the detailed. descriptions found in ?.. for what will henceforth. be called the “UW sample. ?.. for the “Dwarf sample”. and ? for the “ALUGS (MeMaster Unbiased Galaxy Simulations) sample". | We briefly outline the simulations here, but refer the reader to the detailed descriptions found in \citet{Brooks2009}, for what will henceforth be called the “UW sample”, \citet{Governato2010}, for the “Dwarf sample”, and \citet{Stinson2010} for the “MUGS (McMaster Unbiased Galaxy Simulations) sample”. |
ltegardless of the aforementioned "sample! from which a specific simulation was crawn. each individual 7zoomeostyle" simulation was run within a WALAP-8 ACDAL cosmological framework (7). and consisted. of a central high resolution region (centred on the target halo) embedded: within a lower resolution Cosmological volume. | Regardless of the aforementioned `sample' from which a specific simulation was drawn, each individual “zoom-style” simulation was run within a WMAP-3 $\Lambda$ CDM cosmological framework \citep{Spergel2007} and consisted of a central high resolution region (centred on the target halo) embedded within a lower resolution cosmological volume. |
Star formation and supernova feedback was computed using the recipes outlinedby νι with the primary dilference between the three samples | Star formation and supernova feedback was computed using the recipes outlinedby \citet{Stinson2006}, , with the primary difference between the three samples |
The results of the (wo models (thin and thick shell cases) ave shown in Figures 1. and 2.. and summarized in Table 2.. | The results of the two models (thin and thick shell cases) are shown in Figures \ref{fig:rad}~ and \ref{fig:opt}, and summarized in Table \ref{tab:ij}. |
We find that both models provide an equally adequate fit (with Vu71 per degree of freedom). but Ty is not well constrained. Pj~60—5x10* (Figure 3)). | We find that both models provide an equally adequate fit (with $\chi^2_{\rm min}\approx 1$ per degree of freedom), but $\Gamma_0$ is not well constrained, $\Gamma_0\sim 60-5\times 10^3$ (Figure \ref{fig:cont}) ). |
More importantly. both models require collimated ejecta with /;zz0.95 dav (hereafter. Jet). | More importantly, both models require collimated ejecta with $t_j\approx 0.95$ day (hereafter, Jet). |
Models with a significantly wider collimation angle have x5,~LO per degree οἱ freedom. primarily because (μον underestimate (he flux in the X-ray band by a factor of about 20. and cannot explain the radio and optical emission simultaneously. | Models with a significantly wider collimation angle have $\chi^2_{\rm min}\sim 10$ per degree of freedom, primarily because they underestimate the flux in the X-ray band by a factor of about $20$, and cannot explain the radio and optical emission simultaneously. |
At the same time. the Jet models underestimate the radio flux at /Z15 davs by about do on dav 16.3. and about 2.50 on day 32. | At the same time, the Jet models underestimate the radio flux at $t\gtrsim 15$ days by about $4\sigma$ on day 16.3, and about $2.5\sigma$ on day 32. |
This is due to an apparent brightening of the radio⋅ emission⋅⋅ on this ⋅⋅timescale. | This is due to an apparent brightening of the radio emission on this timescale. |
Since⋅ in⋅ the radio⋅ band P,xmj.1/2 one possible. explanation. for the brightening is that the forward shock encounters a density enhancement: a density increase by a factor of ten is required. | Since in the radio band $F_\nu\propto n_0^{1/2}$, one possible explanation for the brightening is that the forward shock encounters a density enhancement; a density increase by a factor of ten is required. |
The flux in the optical bands would remain lareely unallected since Voy.<7. (see Table 2)). in which case the fIuxis independent of density. | The flux in the optical bands would remain largely unaffected since $\nu_{\rm opt}<\nu_c$ (see Table \ref{tab:ij}) ), in which case the fluxis independent of density. |
Using the parameters of the forward shock emission we caleulate Ezz2x10°! ere. e,£Ol eg&0.5. and nyzz107 7 using equations 4.13.4.16 of Sari&Esin(2001). | Using the parameters of the forward shock emission we calculate $E\approx 2\times 10^{51}$ erg, $\epsilon_e\approx 0.1$, $\epsilon_B\approx 0.5$, and $n_0\approx 10^{-2}$ $^{-3}$ using equations 4.13–4.16 of \citet{se01}. |
The opening angle of the jet is 0;zz0.1 (Frailefaf2001). | The opening angle of the jet is $\theta_j\approx 0.1$ \citep{fks+01}. |
. Using this value we lind a beanming corrected 5-rav emergy. £L23.7x10?"50 erg (Priceefal.2002).. twpical for loue- GRBs (Frailεἰal. 2001).. and a beaming corrected kinetic energv. EyzzLO! erg. lower than the tvpical inferred. values of 10?" to 3x107! erg (Panaiteseu&Kumar2002). | Using this value we find a beaming corrected $\gamma$ -ray emergy, $E_\gamma\approx 3.7\times
10^{50}$ erg \citep{pkb+02}, typical for long-duration GRBs \citep{fks+01}, , and a beaming corrected kinetic energy, $E_K\approx
10^{49}$ erg, lower than the typical inferred values of $10^{50}$ to $3\times 10^{51}$ erg \citep{pk02}. |
In the previous section we did not consider the optical emission at /Z10 days since it has a distinct spectrum. Fx4797, compared to Fyx4.PPE! for the early alterglow data. | In the previous section we did not consider the optical emission at $t\gtrsim 10$ days since it has a distinct spectrum, $F_\nu\propto \beta^{-3.9\pm 0.1}$, compared to $F_\nu\propto
\beta^{-1.15\pm 0.07}$ for the early afterglow data. |
Moreover. the predicted brightness of the afterglow at late time is lower by a lactor of about 3—5in the and { bands compared to the flix measured with LIST (Figure 2)). | Moreover, the predicted brightness of the afterglow at late time is lower by a factor of about $3-5$ in the $R$ and $I$ bands compared to the flux measured with HST (Figure \ref{fig:opt}) ). |
These two observations indicate that the late-time emission comes Iron a separate component. | These two observations indicate that the late-time emission comes from a separate component. |
Priceefal.(2002) interpreted this emission as coming from a supernova that occured at about the same time as the burst. | \citet{pkb+02} interpreted this emission as coming from a supernova that occured at about the same time as the burst. |
However. (μον note that the significance of this conclusion depends sensitivelv on the time of the jet break. | However, they note that the significance of this conclusion depends sensitively on the time of the jet break. |
Dased solely on the optical data. these authors were unable to significantly constrain /;. | Based solely on the optical data, these authors were unable to significantly constrain $t_j$ . |
However. our combined radio. optical. and X-ray model with/;%0.95 dav indicates that the SN interpretation is secure. | However, our combined radio, optical, and X-ray model with$t_j\approx 0.95$ day indicates that the SN interpretation is secure. |
We euin further confidence about this interpretation by comparing the Iate-tinme emission | We gain further confidence about this interpretation by comparing the late-time emission |
predict. | predict. |
Saturated conduction is usually defined as heat conduction where electrons carry (heir own thermal energy al their thermal speed. | Saturated conduction is usually defined as heat conduction where electrons carry their own thermal energy at their thermal speed. |
The rationale for (his is that at faster flow speeds (relative to the ions). a Buneman instability would develop. generating Langmuir waves that would inhibit the heat flow. | The rationale for this is that at faster flow speeds (relative to the ions), a Buneman instability would develop, generating Langmuir waves that would inhibit the heat flow. |
Applied to the case above. the saturated heat {lx would be ~3.5x10? eres 7s !. | Applied to the case above, the saturated heat flux would be $\sim 3.5\times 10^5$ ergs $^{-2}$ $^{-1}$. |
In our situation the heat conducting electrons navy. also excite lower-hybrid waves with a similar threshold diit velocity. | In our situation the heat conducting electrons may also excite lower-hybrid waves with a similar threshold drift velocity. |
In fact we should expect the limiting flux to be lower than this simple estimate. as indeed it is. because in our case not only heat conducting electrons but also ions generate the lower-hvbrid turbulence. | In fact we should expect the limiting flux to be lower than this simple estimate, as indeed it is, because in our case not only heat conducting electrons but also ions generate the lower-hybrid turbulence. |
Any electron heating mechanism invoked to explain the observed charge state distributions in the fast solar wind must require an anomalous thermal conductivitv to avoid the deposited heat from being conducted back to the coronal hole. which would contliet with the SUMER temperature ciagnostics of Wilhelmetal.(1998) and Davidetal.(1998). | Any electron heating mechanism invoked to explain the observed charge state distributions in the fast solar wind must require an anomalous thermal conductivity to avoid the deposited heat from being conducted back to the coronal hole, which would conflict with the SUMER temperature diagnostics of \citet{wilhelm98} and \citet{david98}. |
. We consider that the heating and anomalous thermal conductivity provided simultaneously. by. lower-hiybrid waves (o be a desirable aspect of the model. in that one mechanism provides both features. | We consider that the heating and anomalous thermal conductivity provided simultaneously by lower-hybrid waves to be a desirable aspect of the model, in that one mechanism provides both features. |
In (his paper we have argued that a small amount of the energy deposited in ions between 1.5 and 2 R.. should eventually find its wav to the electrons via an instability that generates lower hybrid waves. | In this paper we have argued that a small amount of the energy deposited in ions between 1.5 and 2 $R_{\sun}$ should eventually find its way to the electrons via an instability that generates lower hybrid waves. |
No attempt has been made to address the ion evelotron heating problem. other than to show that the densitv gradients we require are similar {ο those postulated by Markovskii(2001) to generate electrostatic ion cvcelotron waves [from a global resonant. MIID mode. | No attempt has been made to address the ion cyclotron heating problem, other than to show that the density gradients we require are similar to those postulated by \citet{markovskii01} to generate electrostatic ion cyclotron waves from a global resonant MHD mode. |
As such. our line of reasoning is complementary (o that in a recent paper by Cranmer (2003). | As such, our line of reasoning is complementary to that in a recent paper by \citet{cranmer03}. |
. In addressing the larger problem of ion exclotron heating. these authors speculate that low frequency Alfvénn waves can Landau damp on electrons. | In addressing the larger problem of ion cyclotron heating, these authors speculate that low frequency Alfvénn waves can Landau damp on electrons. |
This parallel heating should produce electron beaming and discrete phase-space holes which max heat ions via stochastic processes. | This parallel heating should produce electron beaming and discrete phase-space holes which may heat ions via stochastic processes. |
Thus (he ion heating derives [rom (he electron energization (bv low lrequeney Allvénn waves). rather than the electron heating deriving from the ion heating as in the picture presented here. | Thus the ion heating derives from the electron energization (by low frequency Alfvénn waves), rather than the electron heating deriving from the ion heating as in the picture presented here. |
While a imber of quantitative issues remain unresolved (atomic data. thermal conduction). we believe that an explanation for the observed fast wind elemental charge states in terms of lower hybrid wave electron heating is certainly. plausible. and should be considered. along with other possibilities already discussed in the literature (e.gVocks&Mann2003). | While a number of quantitative issues remain unresolved (atomic data, thermal conduction), we believe that an explanation for the observed fast wind elemental charge states in terms of lower hybrid wave electron heating is certainly plausible, and should be considered along with other possibilities already discussed in the literature \citep[e.g][]{vocks03}. |
.. More accurale numerical models would allow us to exploit this interpretation and allow more rigorous investigation of densitv inhomogeneities in the [ast solar wind. with important consequences for the generation of ion evelotron waves throughout the extended corona. | More accurate numerical models would allow us to exploit this interpretation and allow more rigorous investigation of density inhomogeneities in the fast solar wind, with important consequences for the generation of ion cyclotron waves throughout the extended corona. |
As | As |
Alelotte 66 holds a really special position in our work. | Melotte 66 holds a really special position in our work. |
Whether it should be included in (his case or not is still not certain. | Whether it should be included in this case or not is still not certain. |
Dased on the photoelectric and photographic photometry in [αννοι (1976). the BSSs in Melotte 66 were selected in the cluster central region of 7 arcmin. | Based on the photoelectric and photographic photometry in Hawarden (1976), the BSSs in Melotte 66 were selected in the cluster central region of 7 arcmin. |
However. the BSSs ISED plotted as dotted line in Figure 10 is not at all bluer (han the conventional SSP spectrum. therefore (he B-V color of the cluster is not modified (shown in Table 2). even though it has (he largest Nja οἱ forty-six among all the sample clusters. | However, the BSSs ISED plotted as dotted line in Figure 10 is not at all bluer than the conventional SSP spectrum, therefore the B-V color of the cluster is not modified (shown in Table 2), even though it has the largest $_{BS}$ of forty-six among all the sample clusters. |
Furthermore. we find that forty DS5s of the cluster have rather red spectra instead of blue ones compared. with the turnolf. | Furthermore, we find that forty BSSs of the cluster have rather red spectra instead of blue ones compared with the turnoff. |
The remaining six blue-spectrum BSSs are all located in Ring I of 3 arcmün (Ilwwarden 1976). | The remaining six blue-spectrum BSSs are all located in Ring I of 3 arcmin (Hawarden 1976). |
In. Figure 11. solid circles are the BSSs. | In Figure 11, solid circles are the blue-spectrum BSSs. |
The solid rectangles ave red-spectirum BSSs which should be better classified as vellow stragglers. | The solid rectangles are red-spectrum BSSs which should be better classified as 'yellow stragglers'. |
This complicated situation can be partially attributed to Che significant metallicity dispersion in and around the turnoff region of the cluster (Twarog 1995). apart [rom the membership ancl observational uncertainty issues. | This complicated situation can be partially attributed to the significant metallicity dispersion in and around the turnoff region of the cluster (Twarog Anthony-Twarog 1995), apart from the membership and observational uncertainty issues. |
As one of the most studied old open clusters. NGC 138 has been the subject of numerous studies in many aspects. certainly including5 its BSS population. | As one of the most studied old open clusters, NGC 188 has been the subject of numerous studies in many aspects, certainly including its BSS population. |
All the twenty BSSs in AL95 were selected from Ring5 I and II in the finding5 chart of Sancdage SC(1962). with membership probabilities p > (Dinescu οἱ al. | All the twenty BSSs in AL95 were selected from Ring I and II in the finding chart of Sandage (1962), with membership probabilities p $\geq$ (Dinescu et al. |
1996). | 1996). |
This BSS number almost doubles the number ol eleven from an earlier work of Eggen Sandage (1969). | This BSS number almost doubles the number of eleven from an earlier work of Eggen Sandage (1969). |
AI the BSSs show obvious high central concentration. | All the BSSs show obvious high central concentration. |
In the work of Dinescu et al. ( | In the work of Dinescu et al. ( |
1996). nine out of eleven of the BSSs in Eggen Sanclage (1969) were confirmed as cluster members with high probabilities. while the other two (D and L102 in Sandage 1962) were excluded due to their proper motion membership probabilities. | 1996), nine out of eleven of the BSSs in Eggen Sandage (1969) were confirmed as cluster members with high probabilities, while the other two (D and I-102 in Sandage 1962) were excluded due to their proper motion membership probabilities. |
Dinescu (et al. | Dinescu (et al. |
1996) selected eleven probable BSSs with P, > and D, > (D, ds proper motion membership probability. and D,, is (he combined probabilities of both proper motion and spatial distribution). which include five of Eggen Sandage (1969) while adding six new BSSs. | 1996) selected eleven probable BSSs with $P_{\mu}$ $\ge$ and $P_{\mu,r}$ $\ge$ $P_{\mu,r}$ is proper motion membership probability, and $P_{\mu,r}$ is the combined probabilities of both proper motion and spatial distribution), which include five of Eggen Sandage (1969) while adding six new BSSs. |
Among the six new BSSs. four ave out of Bing-ILI and the other two are in Bine-II of Sandage (1962). | Among the six new BSSs, four are out of Ring-III and the other two are in Ring-II of Sandage (1962). |
Although being the latest observation of the cluster. Dinescus BSS catalog does not show central concentration significantly higher than RGBs. therefore the original AL95 catalog is still adopted in our work. | Although being the latest observation of the cluster, Dinescu's BSS catalog does not show central concentration significantly higher than RGBs, therefore the original AL95 catalog is still adopted in our work. |
The BSS population in this cluster has been followed intensively in the past. | The BSS population in this cluster has been followed intensively in the past. |
Although a few of DS5s are subject to dispute on their membership probabilities. the overall BSs | Although a few of BSSs are subject to dispute on their membership probabilities, the overall BSS |
is AX=3.9L40.73Α. | is $\Delta\lambda=3.94\pm0.73$. |
. Fitting the IL) line of the LRD. we did not get a eood fit. due to the lack of spectral resolution aud the low woad to narrow flux ratio. | Fitting the $\beta$ line of the LRD, we did not get a good fit, due to the lack of spectral resolution and the low broad to narrow flux ratio. |
We could uot fit the bluer Balmer lines because of their low S/N ratios. | We could not fit the bluer Balmer lines because of their low S/N ratios. |
We did fit he Ho line in the new data. | We did fit the $\alpha$ line in the new data. |
Although we do not obtain a eood fit (AZ= L9) because of the complicated line profile tthe two lines lie ou the red aud blue wing of he Πα line). we [Nri]find that there is a broad component with a width of 19A.. while the width of the nebular component is 2.7À. | Although we do not obtain a good fit $\chi^2_{\nu}=4.9$ ) because of the complicated line profile the two ] lines lie on the red and blue wing of the $\alpha$ line), we find that there is a broad component with a width of 19, while the width of the nebular component is 2.7. |
. The [Nit] lines are narrow. with a vpical width of 3A.. quantitatively supporting that the orbidden. nebular lines do not have broad componcuts. | The ] lines are narrow, with a typical width of 3, quantitatively supporting that the forbidden, nebular lines do not have broad components. |
Our new. higlh-resolutiou spectra show narrow nebular lues and broad compouenuts in theHenr. IL. aud Ia lines. | Our new, high-resolution spectra show narrow nebular lines and broad components in the, $\beta$, and $\alpha$ lines. |
Our previous. modoerate-resolutiou spectra show a broad component iu the line. | Our previous, moderate-resolution spectra show a broad component in the line. |
There is no broad component in the nebular lines iu either the new or old spectra. | There is no broad component in the [O ] nebular lines in either the new or old spectra. |
There[O is still no sign of anv absorption lines in the new spectra. | There is still no sign of any absorption lines in the new spectra. |
The broad components of bothTell aud IL have widths ~750 kms. consistent with production in the accretion disk. aud are roughly Cassia. instead of having P-Cweni profiles that would imdicate origin iu a wind. | The broad components of both and $\beta$ have widths $\sim$ 750 km/s, consistent with production in the accretion disk, and are roughly Gaussian, instead of having P-Cygni profiles that would indicate origin in a wind. |
Following Porter(2010).. woe estimate the size of the line-enüttius region. Ay. by assumdüus the lue-cinitting eas is in dKeplerian orbits around a conipaet object. thus Ry,<GAI/0. | Following \citet{porter}, we estimate the size of the line-emitting region, $R_{le}$, by assuming the line-emitting gas is in Keplerian orbits around a compact object, thus $R_{le} \le GM/v^2$. |
We find Ry<2.35vun) AU. which for a mass of 10. AD. would eive an upper limit of 3.1 Rw. | We find $R_{le}<2.35 \left( \frac{M_{\rm{BH}}}{1500\, \rm{M}_{\odot}} \right)$ AU, which for a mass of 10 $_{\odot}$ would give an upper limit of 3.4 $_{\odot}$. |
This is cousistent with oriein of the broad Bine in the accretion disk. | This is consistent with origin of the broad line in the accretion disk. |
The broad line componcuts are shifted relative to th narrow compoucuts. | The broad line components are shifted relative to the narrow components. |
In the new data. the shifts are simall compared to the line width. |56+17 |au/s for and =33420 kms for IL. | In the new data, the shifts are small compared to the line width, $+56 \pm 17$ km/s for and $-33 \pm 20$ km/s for $\beta$. |
These shifts are consistent oulv at the 30 level. which might indicate a difference in the spatial origin of the lines. | These shifts are consistent only at the $3\sigma$ level, which might indicate a difference in the spatial origin of the lines. |
Wowever. this is still consistent with production of both lines within the dis- since random motions within the disk and variatiou between the cussion regions could produce shifts that are siunall compared to the line widths. as observed. | However, this is still consistent with production of both lines within the disk since random motions within the disk and variation between the emission regions could produce shifts that are small compared to the line widths, as observed. |
The ceutral wavelength of theTei broad. component shifts markedly between the odld aud uew data. AA=3.914073 AorAc =252+IF kms. This shift is a substantial fraction of the line width. | The central wavelength of the broad component shifts markedly between the odld and new data, $\Delta\lambda=3.94\pm0.73$ $\,$ or $\Delta v = 252 \pm 47 $ km/s. This shift is a substantial fraction of the line width. |
The shift could be due to random motion within the disk. differing viewing gcomectrics (Robertsotal.2010).. or orbital motion of disk (aud the compact object). | The shift could be due to random motion within the disk, differing viewing geometries \citep{rob10}, or orbital motion of disk (and the compact object). |
If the shifts in the broad component of the liue are due to orbital motion. then this would provide a mcans to deteriuue the orbital period and would also provide a lueasurenmient of the nass fiction for the secondary star. | If the shifts in the broad component of the line are due to orbital motion, then this would provide a means to determine the orbital period and would also provide a measurement of the mass function for the secondary star. |
Thus. a program of niionitoriug NGC 5los N-1 with hieh-resolutiou optical spectroscopic observations will be Huportant iu extending our understanding of the plysical nature of this svsteni. | Thus, a program of monitoring NGC 5408 X-1 with high-resolution optical spectroscopic observations will be important in extending our understanding of the physical nature of this system. |
Iu this section. we make some speculations based on interpretation of the shift in the broad component of the line as due to orbital motion. | In this section, we make some speculations based on interpretation of the shift in the broad component of the line as due to orbital motion. |
One can express the mass function aud the compact object mass. A). dn ternis of the orbital period. P. the velocity excursion. A. aud the companion mass. AL. as where /is the inclination angle aud G is the gravitational constant. | One can express the mass function and the compact object mass, $M_x$, in terms of the orbital period, $P$, the velocity excursion, $K_{x}$, and the companion mass, $M_{c}$ as where $i$ is the inclination angle and $G$ is the gravitational constant. |
From the shift of the line quoted above. we constraiu the scuu-ampltucde of the radial velocity iy,>Δον=126421 lan/s. Thus. if the maxinnun nass of the companion aud the orbital period are kuown. ren Eq. | From the shift of the line quoted above, we constrain the semi-amplitude of the radial velocity $K_x \ge \Delta v/2 = 126 \pm 24$ km/s. Thus, if the maximum mass of the companion and the orbital period are known, then Eq. |
2 leads to an upper bound onthe mass of the conrpact object. | 2 leads to an upper bound onthe mass of the compact object. |
The binary svstem has a visual magnitude ey=22.2 lat eives an upper luit ou the absolute magnitude of i6 colupanion of Vo=6.2 at a distance of L5 Alpe INarachentseyetal. 2002). | The binary system has a visual magnitude $v_0=22.2$ that gives an upper limit on the absolute magnitude of the companion of $V_0=-6.2$ at a distance of 4.8 Mpc \citep{dist}. . |
. Unfortunatelv. this places ittle restriction on the companionmass as even O3V y.ars.With masses of 120. AL... ave allowed. | Unfortunately, this places little restriction on the companionmass as even O3V stars,with masses of 120 $M_{\odot}$ , are allowed. |
However. very | However, very |
histogram clearly slows the coarse iutensitv binning. | histogram clearly shows the coarse intensity binning. |
Without dithering. the image breaks up iuto discrete “bands” of coustaut mutensitv which takes it more difficult to detect faint atures in the nuage. | Without dithering, the image breaks up into discrete “bands” of constant intensity which makes it more difficult to detect faint features in the image. |
This effect is shown in Figure 3.Pa (using aln even nire coarselv quantized Ga = 04 inage to chhance the effect). where the majority of the background pixels all have the same (nediuu erav) intensity value. | This effect is shown in Figure 3, (using an even more coarsely quantized q = 0.5 image to enhance the effect), where the majority of the background pixels all have the same (medium gray) intensity value. |
Iu this first experiment. we exanune how the uncertainties of the photometric aud astrometric nieasurenients of individual stars depend oi the q quantization factor. | In this first experiment, we examine how the uncertainties of the photometric and astrometric measurements of individual stars depend on the q quantization factor. |
To measure this effect. we computed the 5audard Deviation of the magnitude and position errors (ie. the value computed by SExtracOr Wiis he known iuput naenitude or position value) for 6250. «inulaed star Huages at each 0.5. inagnitude lncreiment. | To measure this effect, we computed the Standard Deviation of the magnitude and position errors (i.e., the value computed by SExtractor minus the known input magnitude or position value) for 6250 simulated star images at each 0.5 magnitude increment. |
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