本文提供的数据是Jades Survey38的一部分，特别是通过计划ID 1181获得的（主要研究员D. Eisenstein）获得的 。在两个时期观察到了GN-Z11：UT 5和2023年2月7日的第一个 ，以及2023年5月4日在UT 4和5的第二个
。9。我们参考该论文以详细说明 。简而言之
，使用三个Shutter点点点点击模式，使用NIRSPEC微骨器阵列（MSA）的四种不同配置获得光谱数据。使用四个不同的分散器覆盖0.6–5.3μm波长范围：低分辨率PRISM模式（暴露时间为6,200 s ，每构型）和三个中分辨率的光栅（每个配置3,100 s，每个构型）
，可为统一的未插入插入的Slit39提供R≈1,000的标称光谱分辨率
。但是，GN-Z11的高度紧凑光谱相对于缝隙的宽度 ，导致有效的分辨率更高。为了估计有效的分辨率
，我们将GN-Z11的形态转发为光栅分散器的NIRSPEC仪器，发现分辨率在1,100至2,100之间 。使用了四个MSA配置（两个指向和两个抖动位置）
。可能的观察结果相似 ，但是在这种情况下
，它们由三个连续的dithers（具有三种不同的MSA配置）组成，每个配置都带有三个点头 ，导致棱镜的载载暴露为2.7 h
，三个中等分辨率光栅中的每个均为每个中等分辨率，以及高分辨率的G395H/290LP。不幸的是 ，在MSA上GN-Z11的位置
，后一种光谱在超过4.1μm的波长下被严重截断，因此在此光栅上未观察到所有强的光学发射线 。 　　通过组合两组观测值 ，总暴露时间为9.6 h，而棱镜与6.15 h
，每个中等分辨率光栅
。 　　参考文献中还描述了数据处理。9，我们参考该论文进行详细讨论 。在这里 ，我们只提到我们使用了ESA NIRSPEC科学运营团队和NIRSPEC GTO团队开发的管道
。管道中的大多数处理步骤都采用了JWST Science校准管道中使用的相同算法42。与官方管道不同
，最终的一维（1D）组合光谱是通过组合一维单个光谱而不是在二维光谱中进行提取过程获得的 。此步骤保证了最后的一维光谱对缝隙损失进行了很好的通量校准
。在组合过程中，我们还采用了3σ粘合算法 ，并根据管道提供的数据质量文件排除了坏像素。在科学的基础上
，还优化了单个暴露中1D光谱的提取 。在本文中，我们沿着狭缝采用了三像素（0.3英寸）提取 ，因为它改善了光谱的信噪比（S/N）（用于点源）。最后
，GTO管道通过考虑通量校准处理步骤中的变速箱滤波器吞吐量，为光谱配置G140M/F070LP提供了超出标称波长范围的光谱
。扩展光谱覆盖1.27–1.84μm的波长范围。 　　此外 ，在这里
，我们将光谱光谱结合在其重叠范围中，从而增加了这些地区的S/N 。这些区域中的组合光谱被重新采样至NIV双线周围的8Å ，而不会影响分辨率，而在CIV周围进行了12Å
，因为在这种情况下，连续体需要更高的S/N才能正确地追踪CIV吸收
。 　　在论文中 ，我们采用了来自planck18的扁平λCDM宇宙学
，其H0 = 67.4 km s -1 mpc -1和ωm= 0.315（参考文献43）。 　　使用单个或多个高斯线和用于连续减法的简单功率定律拟合发射线 。使用MCMC（Markov Chain Monte Carlo）算法发现了连续和高斯组件的最佳拟合参数，以估计不确定性
。出于本文的目的 ，除了双线或多重组
，其剩下的紫外线中的每一行都是独立拟合的，除了双线或多重组 ，其线宽度被迫具有相同的值（但请参见下面的讨论）Doublet和多重曲线的相对波长分离
，并将多重次数和多重组的相对波长分开被迫进行名义静电段范围波长。每条线的绝对速度（或在双重组和多重组的情况下）的绝对速度不限制为参考文献中给出的确切红移 。9，允许与快门内目标的位置不确定性相关的小波长校准不确定性
。（目标获取的不确定性可能会导致目标相对于名义位置的偏移约0.05英寸 ，而该管道用于波长解决方案。这种未知的偏移将引入波长的偏移
，最多可通过高达0.5个光谱像素来偏移，该频谱像素对应于不同的范围偏置（给出了不同的范围差异（给出了jevellention distivalent distiveent distipent velements jebibers niffentions） ，该区域是不同的（依赖于依赖的范围） 。交流），还可以在不同的线路之间
，尤其是在BLR中常见的不同线之间的速度偏移
。我们将合适的态度限制在本文的兴趣线上。 　　扩展数据表1提供了拟合的发射线宽度和通量的列表，并提供了扩展数据图1显示了主文本中未显示的附加拟合线 。 　　在NIII]多重的情况下 ，1,748.6Å和1,754.0Å的过渡来自相同的上层
，因此它们的通量比由相关的爱因斯坦系数固定，特别是F1754/F1748 = 1.05。同样 ，1,746.8Å和1,752.2Å来自相同的上层
，其通量比固定为F1746/F1752 = 0.14
。根据GN-Z11光谱推断，推断的线宽度是从线扩散函数反驳的。 　　对于1906年的CIII] ，1906年的双重表
，不可能解决这两个组成部分 。尝试使用两个组件拟合它，使得拟合在两个组件的宽度和强度之间
。此CIII]双重组的其他注意事项是 ，它在正常的星形星系和AGN的NLR中通常也可以看到
，因此它也可以从[NEIII]发射中具有宿主星系的贡献。在参考44，作者使用IFS光谱法来揭示CIII]发射在几个100 PC的尺度上解析 。结果 ，我们在分析中不包括（频谱）未解决的CIII]，因为它不提供对BLR或宿主星系的约束
。在扩展数据表1中
，我们使用单个高斯报告了总通量和宽度。但是，为了完整的目的 ，我们报告说
，通过将两个单独的全宽度组件拟合为最大一半（FWHM）（FWHM），考虑到NLR和BLR ，给出了狭窄组件的C]λ11906/λ11908的c]λ11906/λ11908
，用于fwhms，FWHM的FWHM ，314±120 km s-s-s-s-s-s-s-s-s-s-s-s-s-s-s-s-s-s-s-s-s-s-neiiiiii Iiii Iiii Iiii widths＆560 and and neiii Is）
，和560 si），以及C]λ1908广泛的组件（与NIV宽度一致） 。 　　[OII] 3726,3729 DoubleT可能是在观察到的波长范围内检测到的额外的禁止线 ，可用于约束宿主星系中的速度分散体
。但是
，不幸的是，双线尚未解决。试图拟合它（通过强迫两个组件具有相同的宽度）导致FWHM为365±55 km s -1 ，而通量比为 。 　　在本节中，我们提供有关CIV吸收的一些其他细节
。如文本文所述
，可以看到带有大量Civ BlueShift槽的Civ P-Cygni轮廓与年轻的热星大气有关。然而，此特征的深度是金属性21的强大功能 ，而在GN-Z11中观察到的深槽将需要具有太阳能甚至超值金属度的恒星
，与GN-Z11推断的金属性较低完全不一致 。为了说明与纯恒星起源的不一致性，我们从优雅的HST Survey17中堆叠了11个紫外线光谱 ，并在参考文献中推断出GN-Z11的值周围是金属性。9（z = 0.1 z）
，特别是7.4< 12 + log(O/H) < 7.9. We were conservative by excluding galaxies with strong CIV emission. We also excluded one WR galaxy, as we discuss that the spectrum cannot be dominated by WR stars (see main text and section ‘The WR scenario’). The continuum of the spectra was normalized to one by using a simple linear fit in the spectral ranges 1,410–1,480 Å and 1,560–1,600 Å, consistent with the analysis of the spectrum of GN-z11 in the same spectral region (Fig. 1d). The resulting stacked spectrum is shown with a dashed, orange line in Fig. 1d and illustrates inconsistency with the trough seen in GN-z11. To be conservative, in Fig. 1d, we also show the case of the most extreme spectrum among the 11 selected, the one with the deepest CIV absorption. Although the wings of the stellar winds can extend out to 2,000 km s−1, the profile and depth observed trough at these metallicities is inconsistent with the observed trough in GN-z11 at 5σ. 　　The blueshifted CIV trough (and redshifted emission), therefore, is not a P-Cygni feature associated with stellar (atmosphere) winds. Rather, it is tracing a galactic outflow, as observed in lower redshift starbursts22 and in (mini-)BAL/NAL AGN23,45,46,47,48,49. The determination of the velocity requires knowledge of the exact wavelength of the redshifted, rest-frame CIV transition. Unfortunately, there are small wavelength uncertainties associated with each grating because of the uncertainties of the location of the sources within the shutter, as discussed above. In this specific case, we calibrate the velocity shift based on the NIV line, which is in the same gratings and has a similar ionization potential as CIV. The outflow velocity is also subject to different definitions. The centroid of the trough relative to the mean of the two CIV transitions gives a velocity of –790 km s−1. If we consider the blue edge of the trough relative to the bluest of the two transitions (CIVλ1548.19), then we obtain a velocity of –1,040 km s−1. These velocities are significantly higher than those inferred from the CIV absorption in starburst-driven outflows22,50, but in the range of BAL quasars that can span from 500 km s−1 to several thousands km s−1 (refs. 23,45,46,47,48,49,51,52). 　　The classification boundary between mini-BAL and NAL AGN is not sharp, with different authors giving different definitions in terms of width and/or blueshift of the absorption23,45,46,47,48,49. Here we simply give a generic classification as mini-BAL/NAL without aiming at a more specific category. 　　It should be noted that some past works have reported some rare starburst galaxies showing outflows with high velocities, even in excess of 1,000 km s−1 (refs. 53,54,55,56). However, these outflows are traced by lower ionization transitions (MgII absorption and [OIII] emission). More importantly, a close inspection of those cases reveals that each of them shows some AGN signature ([NeV] emission and/or broad MgII emission and/or broad Hβ emission and/or X-ray emission and/or located in the AGN or composite region of diagnostic diagrams). Therefore, although the AGN contribution to the bolometric luminosity of these galaxies may be arguable (also taking into account the variable nature of AGN), it is likely that the high-velocity outflows seen in these rare cases are actually driven by the AGN that they host. 　　Finally, it should be noted that the CIV absorption trough goes nearly to zero (as in many BAL quasars), which implies the total covering factor of the emitting source by the outflowing ionized gas along our line of sight. However, the errors leave scope for a contribution of 30% of the emission potentially not covered by the CIV absorption, which can be associated with the extended host galaxy. Yet, if higher S/N data confirm the CIV trough going to zero, this would imply that the outflow has an extent covering also the host galaxy, that is, about 400 pc, which would be fully consistent with recent findings of BAL outflows extending on scales of up to several kpc (refs. 49,57,58,59). 　　Given that CIV is a resonant line, the observed redshifted emission is also tracing the CIV counterpart of the redshifted Lyα emission seen in ref. 9—that is, the receding side of the outflow. 　　We finally note that the spectrum between Lyα and the NV doublet shows the tentative signature of an NV blueshifted trough (Extended Data Fig. 1d), which would be associated with the highly ionized outflow, but it requires additional data to be confirmed. 　　Although the paper focuses on a few lines discussed in the main text, in this section we also discuss other emission lines that have either lower S/N, more severe blending or whose upper (or lower) limits give line ratios that are fully consistent with the AGN scenario. 　　The MgII2796,2804 doublet is well resolved with the grating and in principle a good tracer of gas density in the range between 109 cm−3 and 1014 cm−3. However, the observed ratio, , is so uncertain to be consistent both with the low-density regime (ratio of around 1) and the high-density regime (ratio of about 2). Moreover, even if additional data allows constraining the MgII doublet ratio more tightly, these are resonant transitions, which are, therefore, strongly sensitive to the optical depth and radiative transfer effects60. 　　The C]λλ1907,1909 doublet would also be a good density tracer, as the ratio of its two components is primarily sensitive to the gas density and changes strongly between 104 cm−3 and 106 cm−3 (with the blue component λ1907 going to zero at high densities), similar to the NIV] doublet. However, as discussed above, the two components are unresolved with the grating, and we cannot obtain reliable constraints on the gas density or on the line widths. More importantly, CIII] emission is commonly seen also in star-forming galaxies and in the NLR of AGN, so it may partially come also from the low-density ISM of the host galaxy, as is the case for [NeIII]. As already mentioned, recent IFS observations show CIII] to be resolved on scales of several 100 pc (ref. 44). It is interesting that when fitted with narrow and broad components, as discussed in the previous section, the narrow component gives widths formally consistent with the [NeIII], whereas the broad component is consistent with the NIV width. 　　Additional transitions from species requiring ionizing photon energy higher than about 60 eV, such as NV and NeV (in addition to NeIV seen in GN-z11), are often seen as evidence for the presence of an AGN. Yet, conversely, their absence should not be necessarily seen as evidence for the absence of an AGN, as often these lines are weak even in AGN and remain undetected if the S/N is not high enough61,62,63. Moreover, the intensity of these lines varies strongly from case to case. 　　With the prism it is not possible to assess the presence of NV because it is blended with Lyα and its damping wing. Regarding the gratings, the G140M band, in which NV is redshifted, is the least sensitive of the three medium-resolution spectra. Although there is a hint of the NV doublet (Extended Data Fig. 1a 2σ integrated signal) we obviously do not quote it as a tentative detection. The inferred upper limit on the NV emission is not very constraining, but the important aspect in the context of this paper is that it is still fully consistent with the presence of an AGN. We demonstrate this in Extended Data Fig. 2, in which the upper limits on the NV/CIV and NV/HeII ratios for GN-z11 are compared with a sample of the broad lines in type 1 AGN61 and also with a sample of the NLR in type 2 AGN62, and illustrating that the non-detection of NV is fully consistent with the AGN scenario. 　　It is also interesting to compare NV with NIV, as this ratio is not dependent on the nitrogen abundance, although NIV is detected (or reported) less frequently in AGN. In the well-studied type 1.8 AGN at z = 5.5, GS-3073 (refs. 16,30,64), the NV is five times fainter than NIV, which would be totally undetected in our spectrum. In the type 1 quasars explored in ref. 65, the NIV broad line is very strong, whereas NV is undetected, with an upper limit that is about 10 times lower than the NIV flux. 　　NeV is also not detected, neither in the grating nor in the prism spectrum. The upper limit on the NeV/NeIII ratio is about 0.2. However, in ref. 63, the authors have shown that AGN models can have NeV/NeIII as low as 10−2−10−4. Hence the non-detection of NeV is also not constraining about the presence of an AGN. 　　Finally, we note that AGN accreting at super-Eddington rates have a lower energy cutoff, and hence are less likely to emit hard photons that can produce highly ionized species, such as NV and NeV. 　　HeII is detected in the prism and, more marginally, in the grating (Extended Data Fig. 1). 　　As already discussed, CIV is detected in the grating, but with a P-Cygni profile, hence its flux is a lower limit because of self-absorption. 　　The interpretation of these limits using photoionization models is very much model-dependent. We illustrate this in Extended Data Fig. 3a,b. Specifically, Extended Data Fig. 3a, as in ref. 9, shows the location of GN-z11 on the CIII]/CIV versus HeII/CIII] diagram and in which the red-squared and blue-starred symbols show the location of models from refs. 10,66 for the NLR of AGNs and for star-forming galaxies, respectively, and in a range of about ±0.3 (see legend) dex of the metallicity inferred in ref. 9 for GN-z11. GN-z11 can be consistent with both AGN and star-forming models. 　　Extended Data Fig. 3b shows the same diagram in which we instead plot the models from ref. 67, in the same (low) metallicity range for both AGN and SF galaxies. In this case, GN-z11 is much more consistent with the AGN models and inconsistent with the models for star-forming galaxies. 　　Yet, if the permitted and semi-forbidden lines of GN-z11 are coming from the BLR, as argued in this paper, then neither of the models above actually apply, as they are developed for the low-density environments of the NLR and HII regions. It is, therefore, more instructive to compare with the line ratios observed in the BLR of type 1 AGN. These are taken from the compilation of ref. 61 and shown with purple circles in Extended Data Fig. 3c. The line ratios observed in GN-z11 are fully consistent with the broad lines of type 1 AGN. For completeness, in the same panel we also plot the ratios observed for the NLR of type 2 AGN, compiled in ref. 62 (mostly overlapping with the ratios observed for the broad lines), and the star-forming galaxies from the CLASSY survey18. 　　In this section, we discuss the scenario recently proposed that GN-z11 may be similar to local WR galaxies24. 　　The HeII marginal detection shows a potentially broad profile (about 10,00 km s−1, although the broad wings are mostly in the noise), which may be associated with the inner BLR, but also may resemble the broad HeII profile characteristic of WR stars. Therefore, there might be a contribution from WR stars and, specifically, WN stars, given the strong nitrogen lines. 　　However, there are various spectral features that cannot be accounted for in the WN scenario. 　　WN stars are also characterized by very strong NIVλ1718 resonant emission, stronger than the NIVλ1486, and typically with a prominent P-Cygni profile13. In GN-z11, despite the very strong NIVλ1486, there is no trace of the NIVλ1718 line. Figure 1f shows the spectrum of GN-z11 at the expected location of NIVλ1718 and in which the shaded red region shows the expected intensity of the line, based on the strength of the NIVλ1486 line. The GN-z11 spectrum is totally inconsistent with the presence of the NIVλ1718 WR signature. 　　Furthermore, neither [Ne]λ2424 nor CII*λ1335 are ever seen associated with the WR population13. 　　Finally, even if WN show prominent NIII] emission, the strength of the λ1754 component of the multiplet is much fainter in WR galaxies such as Mrk966 (ref. 17) and consistent with densities typical of the ISM. 　　In sum, although WR stars might be present in GN-z11, they are unlikely to dominate the excitation of most nebular lines. 　　Extended Data Table 2 summarizes more schematically the observational features consistent or inconsistent with the AGN scenario, the WR scenario and a compact starburst without WR stars. 　　We used the C photoionization code68 to explore the effect of varying physical conditions on some emission line ratios constrained by JWST/NIRSpec. The primary goal is to explore the ratios of emission lines within a given doublet or multiplet, hence lines of the same ion (specifically NIII and NIV) that are effectively insensitive to the chemical abundance and ionization parameter, while sensitive to density and only with secondary dependence on temperature. For this reason, the details of the photoionization models are not as important as when exploring other line ratios. We considered a nebula of constant pressure in plane-parallel geometry. However, we have verified that other scenarios, such as a cloud with constant density, do not affect our findings. For completeness, we considered both AGN and stellar templates for the shape of the incident radiation field. Its normalization is set by the ionization parameter, defined as U ≡ ΦH/(nHc), where ΦH is the surface flux of hydrogen-ionizing photons at the illuminated face of the nebula, nH is the number density of hydrogen and c is the speed of light. The hydrogen density and ionization parameter were varied in logarithmic steps of 1, respectively from nH = 1 cm−3 up to nH = 1014 cm−3, and starting at log10U = −3 and ending at log10U = −1 (refs. 10,66,67). 　　In the AGN scenario, we adopted the multi-component continuum template implemented in C, consisting of a black body and a power law, varying the black body temperature (TAGN = 106 K and 106 K) while fixing the power-law slope to α = −1.4 (note that this is the slope underlying the black body at energies above the Ly-edge) and leaving other optional parameters as default. For the AGN models, we considered gas-phase metallicities of Zneb = 0.1 Z and Zneb = 1 Z. By contrast, the star-forming models are restricted to Zneb = 0.1 Z, as the hard ionizing spectra of metal-poor stars are essential to form sufficient triply ionized nitrogen (requiring 47.5 eV), whose presence in GN-z11 is evidenced by the strong NIV emission (EWNIV 1486 = 9.0 ± 1.1 Å; ref. 9), whereas metal-rich stars would not produce enough hard ionizing photons to make the NIV line visible. In the star-formation scenario, we used stellar population synthesis models, including binary stars generated by v.2.1 (ref. 69) for a single burst of star formation (with varying ages, t*/Myr  {1, 10, 100}), assuming the same metallicity as the gas (that is, Z* = Zneb = 0.1Z) and an IMF70 that ranges in stellar mass from 1M to 100M. Both in the AGN and star-formation cases, calculations are run until a neutral hydrogen column density of NHI = 1021 cm−2 is reached to ensure that in all models the nebula is matter bounded; we note, however, that the highly ionized nitrogen lines are produced in the very inner part of the cloud, such that the boundary conditions do not significantly affect our results. In total, this results in a parameter grid of 15 different densities, 3 ionization parameters, 3 temperatures or stellar ages, 2 or 1 metallicities for the AGN and star-formation models, respectively, or a total of 15 × 3 × (3 × 2 + 3 × 1) = 405 possible model configurations. 　　The relevant nitrogen line ratios for all of these (except for eight cases in which C reported a failure) are shown in Fig. 2, from which we conclude that they are consistent between the AGN and star-formation scenario, and their density dependence is largely independent of ionization parameter, metallicity or the precise shape of the incident radiation field (that is, AGN or star formation and the corresponding parameter TAGN or t*). 　　At high densities, the NIVλ1483/NIVλ1486 ratio approaches zero (nH  106 cm−3), whereas N]λ1754 plateaus at a fractional contribution to the multiplet of about 0.23 at higher densities still (nH  1010 cm−3), both pointing towards the presence of a broad line region in GN-z11 given the observed values. 　　Finally, to increase the readability of Fig. 2, we have separated the AGN and star-forming models in two separate panels in Extended Data Fig. 4. 　　If GN-z11 is a type 1 AGN, then we should be directly seeing the light from the accretion disc. In the case that the accretion disc dominates, the UV-to-optical continuum should follow a simple power law of the form Fλ  λβ with β = −7/3 ≈ −2.33 (ref. 71), as observed in type 1 AGN, and NLSy172,73, modulo the UV turnover whose wavelength increases with black hole mass and also modulo effects of dust reddening, which often makes the spectrum redder. 　　In the case of GN-z11, the spectrum is contributed to also by the underlying galaxy identified in ref. 8 in the NIRCam images. This component is significantly fainter than the nuclear point-like component. It is difficult to quantitatively establish its contribution to the spectrum, because part of the light may fall outside the shutter, and in a different fraction in the four dither or pointing positions, and not easy to reconstruct because of the slight positional uncertainties discussed above. In Extended Data Fig. 5a, we show the contribution from the galactic component (dotted-orange line) to the spectrum, assuming that the entire light of the galaxy is captured by the spectrum, corresponding to about one-third of the flux and using the spectral template inferred in ref. 8 for the extended component. 　　The additional component to take into account is the nebular continuum associated with the BLR (as well as any other ionized gas in the host galaxy). The BLR typically has a low covering factor74, therefore the nebular continuum is not expected to be strong, but its contribution must be quantified. In most physical conditions typical of the ionized gas in the BLR, NLR or HII regions, the nebular continuum is linked to the intensity of the Balmer lines. We have estimated the nebular emission using a Cloudy model with a metallicity of 0.1 Z and a density of 106 cm−3 (between the BLR and ISM origin scenarios) and normalized to have the same Hγ flux as observed in the spectrum of GN-z11. The nebular spectrum does not change drastically as a function of density, except obviously for the emission of the forbidden and semi-forbidden lines; however, our focus is on the nebular continuum, so we ignore the mismatch of the emission lines, as a detailed photoionization modelling of their flux is beyond the scope of this paper. We note that the nebular continuum is also included in the model spectrum fit to the extended component in ref. 8; therefore, not to include it twice, we have measured the Hγ flux in the ref. 8 spectrum and normalized the Cloudy nebular spectrum only to the Hγ flux obtained by the difference between the observed value and the flux in the ref. 8 model spectrum. The resulting nebular spectrum is shown with a dashed purple line in Extended Data Fig. 5a. Again, the mismatch of the emission lines should be disregarded, as the goal is not to reproduce them with the Cloudy model. 　　Extended Data Fig. 5b shows again the observed spectrum, in log–log scale, in which the galactic and nebular components have been subtracted. Although the noise is large, especially at long wavelengths, also as a consequence of the model subtraction procedure, the resulting spectrum is well fitted by a simple power law, in the parts not affected by the emission lines. The best-fitting slope is –2.26 ± 0.10, hence consistent with the continuum expected from an accretion disc. Note that this is not evidence in support of the presence of an AGN, as also young galaxies may have power-law shapes, it is only meant to show consistency with the AGN scenario. 　　We finally note that, although with a large scatter, the UV spectrum of AGN often shows a FeII hump between about 2,300 Å and about 3,100 Å (refs. 75,76,77,78,79). The prism spectrum of GN-z11 does not show an obvious FeII bump, although a more detailed analysis and modelling is required to assess the presence or absence of such a bump, which is deferred to a separate paper. However, we note that at such early epochs there is little time for the ISM to be enriched with iron through the SNIa channel80, so a weak or absent FeII bump would not be unexpected. 　　The luminosity of AGN can be variable, from a few per cent to a factor of a few, on short (days) and long (years) timescales. We have investigated the possible presence of variability. Before the recent NIRCam images obtained in February 2023 (ref. 8), deep photometric observations were obtained with HST about 10 years earlier81,82, corresponding to about 1 year in the rest frame of GN-z11. Most of the HST photometric data points have error bars that are too large to be useful for constraining variability. However, the photometric point reported in ref. 81 with the F160W filter has a relatively well constrained value: 150 ± 10 nJy, within an aperture of 0.35″. NIRCam does not have the same filter, however, the photometry obtained in the F150W filter can be used and transposed to the F160W filter by using the NIRSpec prism spectrum. After extracting photometry from a 0.35″ aperture (as in ref. 81), and extrapolating with the NIRSpec spectrum, we obtain a F160W equivalent photometry of 141 ± 2 nJy, which is consistent with the HST previous photometry within 1σ. If we consider that about 30% of the flux is diluted by the host galaxy, the comparison of the photometry between the two epochs would indicate a variability of 10% at only 1σ. This is certainly not a detection of variability, but it is consistent with the range of variability amplitudes observed in NLSy1 and, more broadly, in type 1 AGN83. 　　GN-z11 is not detected in X-rays. Flux limits are obtained from the Chandra Deep Field North, which was a 2 Ms observation performed in 2002 (see84 for final results). Their sensitivity map gives a point source limit in the soft (0.5–2 keV), hard (2–7 keV) and full (0.5–7 keV) bands of 1.54 × 10−17, 7.9 × 10−17 erg cm−2 s−1 and 4.9 × 10−17 erg cm−2 s−1. Source detection requires a no-source probability P < 0.004. The tightest limit in the soft band translates to a rest frame 5.8–23.2 keV luminosity limit at z = 10.6 of 2.2 × 1043 erg s−1. Assuming a typical NLS1 photon index of 2.3 means that LX (2–10 keV) is less than 3 × 1043 erg s−1. 　　The bolometric correction for NLS1 in the 2–10 keV band, BCX, is about 100 (ref. 85). There is a significant systematic uncertainty here due to the unseen flux in the FUV, in which the emission is expected to peak (see fig. 3 in ref. 86). Moreover, the 2–10 keV flux entirely originates from the corona, the early development of which and possible dependence on black hole spin are unknown (ref. 87 cautions against using his X-ray BC values for NLS1). Proceeding with BCX = 100 means that the Chandra upper limit is almost three times above the luminosity inferred from the JWST flux at 1,400 Å. We predict a conservative SB flux of 5 × 10−18 erg cm−2 s−1. This would be detectable in about 1 Ms with the candidate NASA Probe mission AXIS. The coronal emission from local NLS1s is highly variable and the above BC represents a mean value (note that the intrinsic disc flux seen in the UV is much less variable86). 　　For the vast majority of high redshift AGN, the black hole masses are inferred using single-epoch measurements and the so-called virial relations, that is, relations between the black hole mass, the width of the lines of the BLR and the continuum or line luminosity37,88,89,90,91,92,93,94. These relations are calibrated on nearby AGN, using either reverberation mapping techniques and/or direct dynamical measurements of the black hole. The black hole mass scales about as the square power of the width of the BLR lines and about as the square root power of the luminosity, with a proportionality constant that depends on the specific waveband (or line) for the luminosity estimation. 　　The most accurate virial relations would be those using Hα and Hβ. In our case, Hγ could be used as a proxy. However, as discussed, the Balmer lines are probably contributed to by the star formation in the host galaxy, hence not reliable to trace the black hole mass. 　　The CIII] doublet is also sometimes used to infer the black hole mass. However, this is not well resolved and, as for the case of the Balmer lines, this is probably contaminated by the ISM and star formation in the host galaxy. 　　MgII is often used. In our case, the MgII doublet is clearly detected, but the S/N is fairly low for the measurement of the width (Extended Data Fig. 1). If we take the width resulting from the fit and the relation provided in ref. 95: 　　then we get a black hole mass of 1.4 × 106M. However, given the low S/N on the MgII doublet, we prefer to use as representative width of the BLR lines the profile of the high S/N and isolated NIV line. If we adopt this width into the equation above, we obtain a black hole mass of 1.6 × 106M. The uncertainty is totally dominated by the scatter in the virial scaling relation, which is about 0.3 dex (ref. 96). 　　Moreover, there are various other systematic uncertainties and caveats that can affect the black hole mass estimate. To begin with, it is not obvious that the local virial relations apply at high redshift. The main issue is whether the dependence of the BLR radius on luminosity evolves with redshift or not. The most plausible scenario is that the square root dependence of the BLR radius from luminosity is primarily set by the dust sublimation radius. In ref. 100,101) are shown in Fig. 4. These are based on the Hα or Hβ width and flux. We clarify that these are re-estimated by using the same calibrations used in ref. 36 for local galaxies. 　　We derive the bolometric luminosity of the AGN by using the continuum luminosity at λrest = 1,400 Å and the luminosity-dependent bolometric correction given in ref. 87: 　　We also assume, as discussed in the previous sections, that 30% of the continuum flux at this wavelength is because of the underlying galactic component8 and that, therefore, the AGN continuum luminosity at this wavelength is 0.7 of the observed value. We infer a bolometric luminosity of 1.08 × 1045 erg s−1. The resulting ratio between bolometric and Eddington luminosity is 5.5, also affected by an uncertainty of a factor of at least 2, coming from the uncertainty on the black hole mass. 　　There is a vast literature discussing the formation of early black holes and on how they evolve in the first thousand million years, by using hydrodynamical and cosmological simulations, as well as semi-analytical models. The production and elaboration of models in this area have recently seen surge with the goal of specifically interpreting the results from JWST. It is beyond the scope of this paper to provide an exhaustive description of the assumptions and results of the several models and simulations. However, in this section, we briefly discuss that many of them can explain the properties of GN-z11 and provide some possible constraints on the seeding scenarios. 　　We start by considering the results obtained in ref. 7 from the FABLE hydrodynamical, cosmological simulation, in which they focused on the largest halo at z = 6 (with a virial mass M200 = 6.9 × 1012M of the Millennium box). The latter may appear an extreme choice; however, we note that GN-z11 does live in an overdense region and probably at the core of a protocluster8,102. In the FABLE simulation, the black hole seed has a mass of 105M at z = 13. The accretion rate is capped to Eddington and uses the Bondi–Hoyle–Littleton-based formalism; however, as small scale, non-isotropic accretion is unresolved in the simulation, FABLE, like Illustris, uses a Bondi–Hoyle–Littleton rate boosted by a factor of 100. Feedback energy in FABLE scales as 10% of the available accretion energy, , where ϵ = 0.1 is the radiative efficiency of the accretion flow. At high redshifts, this is primarily injected as thermal energy in the vicinity of the black hole, with a duty cycle of 25 Myr. We overplot the fiducial model in ref. 7 in Fig. 3 (orange solid line, labelled as B23), illustrating that this can easily reproduce the mass of the black hole in GN-z11 at z = 10.6. 　　To explain the most massive BHs observed at z ≈ 6–7, the same study as above 7 also explores the scenario of earlier seeding (z = 18) and allows the black hole to accrete at up to two times the Eddington limit; in this case, the model could explain a black hole nearly five times more massive than GN-z11 at z = 10.6. 　　In ref. 35, the authors explored the early evolution of black holes using the TRINITY cosmological empirical model103, which is based on halo statistics from N-body simulations and incorporating empirical galactic scaling relations. The authors specifically explore the case of GN-z11. They illustrate that its mass and black hole to stellar mass ratio can be explained by their model starting with an intermediate mass seed of a few times 103 seeded at z = 15, accreting on average at sub-Eddington rates, but intermittently also at super-Eddington. Their track is shown with a solid-teal line in Fig. 3 (labelled as Z23). 　　Recently, in ref. 25, the authors have explored the properties of GN-z11 within the context of the semi-analytical model CAT. They find that the black hole mass of GN-z11 and its location on the MBH–Mstar diagram can be interpreted both in terms of light seeds (at z = 20–23) that can have super-Eddington accretion phases, or Eddington-limited heavy seeds formed at z = 14–16. Out of their various tracks, Fig. 3 shows only two samples of their tracks, in the case of a light (red-solid) and a heavy seed (red-dashed), which can both reproduce the mass of GN-z11 at z = 10.6 (labelled as S23). In both cases, the semi-analytical model can also reproduce the black hole to stellar mass observed in GN-z11. 　　In ref.  9 JWST galaxies hosted massive black holes? Mon. Not. R. Astron. Soc. 521, 241–250 (2023)." href="https://www.nature.com/articles/s41586-024-07052-5#ref-CR3" id="ref-link-section-d30367176e4372">3, the authors suggested that the detectability of accreting BHs at high redshift by JWST implies that these are probably originating from heavy seeds. Specifically, their models can reproduce the mass of GN-z11 at z = 10.6 but only with seeds that are several times 105M, already in place before z = 14. GN-z11 would fall in this category, and the tracks obtained in ref.  9 JWST galaxies hosted massive black holes? Mon. Not. R. Astron. Soc. 521, 241–250 (2023)." href="https://www.nature.com/articles/s41586-024-07052-5#ref-CR3" id="ref-link-section-d30367176e4388">3还将解释在GN-Z11中观察到的黑洞与恒星质量比
。 　　其他研究提出了其他方案，可视化不同的播种机制 ，在不同的红移处
，并具有对积聚和合并速率的不同假设，并且能够以z = 10.6的gn-Z11的黑洞质量再现黑洞的黑洞质量 ，并且通常也可以将其黑洞与恒星质量比率为4,26,27,27,27,104。 　　总而言之
，可以使用不同的假设来解释GN-Z11中黑洞的性能，这些假设可以在以亚埃德丁顿速率或中间的灯光中的重型种子进行大致分组 ，或者以超级 - 埃德丁顿相和/或以增强的Bondi积聚为模型 。 　　需要更多关于GN-Z11的物体的统计信息才能区分不同情况
。目前，据我们所知
，GN-Z11仍然是所有HST深区（包括蜡烛和边境领域）中Z> 10的最发光物体。希望JWST对较大区域（例如，在Cosmos-Web中）的观察能够在z> 10上找到类似于GN-Z11的AGN 。目前 ，正如文本中所讨论的那样
，有趣的是，模型和模拟期望有一些黑洞 ，质量为106-107m
，在10< z < 11 in the JADES Medium-Deep survey in the GOODS fields 9 JWST galaxies hosted massive black holes? Mon. Not. R. Astron. Soc. 521, 241–250 (2023)." href="https://www.nature.com/articles/s41586-024-07052-5#ref-CR3" id="ref-link-section-d30367176e4434">3,105
。因此，在GN-Z11中发现2×106m黑洞并不意外 ，并且在货物场中可能会出现更多（可能以较低的速度吸收）。 　　我们已经表明
，GN-Z11的高氮富集可能仅限于BLR，BLR的质量小和紧凑的大小可能经历了非常快速的化学富集 ，只需要几个SNE 。 　　我们注意到
，GN-Z11的高化学富集与最近在GN-Z1144光环中对原始气体的主张不符
。这些主张在完全不同的尺度上，原始气体与GN-Z11相距几个KPC ，而高化学富集估计为GN-Z11的核。关于原始气体在GN-Z11的光环中的主张，模型预计高Z的大型星系可能会在光环中容纳原始气体的口袋
，甚至降至Z≈3（参考文献106,107） 。
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